diff --git "a/9802.json" "b/9802.json" new file mode 100644--- /dev/null +++ "b/9802.json" @@ -0,0 +1,1607 @@ +{ + "9802/astro-ph9802239_arXiv.txt": { + "abstract": "We discuss the rotational excitation of small interstellar grains and the resulting electric dipole radiation from spinning dust. Attention is given to excitation and damping of grain rotation by: collisions with neutrals; collisions with ions; ``plasma drag''; emission of infrared radiation; emission of electric dipole radiation; photoelectric emission; and formation of $\\HH$ on the grain surface. Electrostatic ``focussing'' can substantially enhance the rate of rotational excitation of grains colliding with ions. Under some conditions, ``plasma drag'' -- due to interaction of the electric dipole moment of the grain with the electric field produced by passing ions -- dominates both rotational damping and rotational excitation. We introduce dimensionless functions $F$ and $G$ which allow direct comparison of the contributions of different mechanisms to rotational drag and excitation. Emissivities are estimated for dust in different phases of the interstellar medium, including diffuse HI clouds, warm HI, low-density photoionized gas, and cold molecular gas. Spinning dust grains could explain much, and perhaps all, of the 14 - 50 GHz background component recently observed by Kogut et al. (1996), de Oliveira-Costa et al. (1997) and Leitch et al. (1997). Future sensitive measurements of angular structure in the microwave sky brightness from the ground and from space should detect this emission from high-latitude HI clouds. It should be possible to detect rotational emission from small grains by ground-based pointed observations of molecular clouds. ", + "introduction": "Experiments to map the cosmic microwave background radiation have stimulated renewed interest in diffuse Galactic emission. Most recently, Kogut et al (1996), de Oliveira-Costa et al. (1997) and Leitch et al. (1997) have reported a new component of galactic microwave emission which is correlated with $100\\micron$ thermal emission from interstellar dust. Kogut et al. found the emission excess to have $I_\\nu(31.5{\\rm GHz})\\approx I_\\nu(53{\\rm GHz})$, and Leitch et al. found $I_\\nu(14.5{\\rm GHz})\\approx I_\\nu(32{\\rm GHz})$, consistent with the spectrum of free-free emission, but nondetection of H$\\alpha$ emission in these directions is inconsistent with free-free emission accounting for the microwave excess unless the plasma temperature $T\\gtsim 10^6\\K$. Leitch et al. therefore proposed that the observed emission was free-free emission from shock-heated gas in a supernova remnant. Draine \\& Lazarian (1998, hereafter DL98) showed, however, that the observed microwave excess could not be due to free-free emission from hot gas, as this would require an energy injection rate at least 2 orders of magnitude greater than that the energy input due to supernovae. DL98 showed that the microwave excess could in fact be electric dipole emission from rapidly-rotating dust grains. To predict the intensity of dipole emission one needs (1) the numbers of small grains, (2) their dipole moments, and (3) their rotational velocities. The observed intensity of 12 and $25\\micron$ emission from interstellar clouds allows us to estimate the numbers of very small grains (Leger \\& Puget 1984; Draine \\& Anderson 1985; Desert, Boulanger, \\& Puget 1990). To estimate the dipole moments, we consider the likely displacements between the charge and mass centroids for grains, plus the intrinsic dipole moment arising from polarized chemical bonds in the grain material. However, the main thrust of the present paper is a comprehensive study of the rotational dynamics of small grains, in order to estimate their rotation rates. Ferrara \\& Dettmar (1994) showed that small grains rotating with thermal rotation rates could produce detectable radio emission, but they did not address the details of rotational excitation and damping. Rotational excitation of small grains (= large molecules) has been discussed previously by Rouan et al (1992), who described the effects of collisions with gas atoms and absorption and emission of radiation. In the present work we reexamine this problem, and include the important effects of collisions with ions, which were neglected in the study by Rouan et al. We also include the effects of ``plasma drag'' due to interaction of the grain with passing ions, first considered by Anderson \\& Watson (1993). We derive rates for rotational damping and excitation due to plasma drag which are somewhat larger than estimated by Anderson \\& Watson (1993). We discuss rotational excitation as a function of both grain size and environmental conditions. Electrostatic focussing of ions makes them very effective at delivering angular momentum to the grains; ion collisions can dominate the rotational excitation of very small grains even in predominantly neutral regions. For our adopted population of very small grains, we predict the microwave emissivity of various phases of the interstellar medium, ranging from diffuse gas to molecular clouds. We expect detectable levels of emission from spinning dust grains in all of these phases. The paper is organized as follows: After describing the environments which we consider (\\S\\ref{sec:media}), we discuss the electric dipole moments expected for small dust grains (\\S\\ref{sec:grain_properties}). The various rotational damping processes are reviewed in \\S\\ref{sec:rotational_damping}, and the rotational excitation mechanisms in \\S\\ref{sec:rotational_excitation}. Using these rates for rotational excitation and damping, we calculate the resulting rate of grain rotation in \\S\\ref{sec:rot_rate}. The importance of the impulsive nature of the rotational excitation is discussed in \\S\\ref{sec:impulsive_J}, and the effects of centrifugal stresses in \\S\\ref{sec:centrifugal}. The reader interested primarily in the predicted emission may wish to skip \\S\\S\\ref{sec:media}--\\ref{sec:centrifugal} and proceed directly to \\S\\ref{sec:sized}, where we describe our assumptions concerning the grain size distribution in various environments. We present the microwave emission spectra expected for grains in both diffuse regions and molecular clouds in \\S\\ref{sec:emissivity}. The detection of this microwave emission from pointed observations of dense clouds is discussed in \\S\\ref{sec:darkcloud}. The principal uncertainties in our estimates are discussed in \\S\\ref{sec:discussion}, and our results are summarized in \\S\\ref{sec:summary}. ", + "conclusions": "} We have presented above a detailed study of rotational excitation and damping for grains with sizes less than $\\sim10^{-6}\\cm$. In addition to processes included in earlier work by Rouan et al. (1992), we have included both direct collisions with ions and ``plasma drag''. Ion collisions and plasma drag, together with damping by infrared and microwave emission, dominate rotational excitation and damping for most interstellar environments. It is therefore important to include these processes in other studies of rotational excitation, such as those proposing to explain features of the diffuse interstellar bands (Rouan, Leger, \\& Coupanec 1997). For the assumed distribution of grain sizes and electric dipole moments, we predict microwave emission from spinning dust grains which can account for the ``anomalous emission'' observed recently in the 15--90 GHz range. The predicted intensities are uncertain, however, as they depend on three poorly-known factors: \\begin{enumerate} \\item The abundances and size distribution of the small grains. This uncertainty is greatest in the case of dense/dark regions, where the absence of starlight denies us the evidence (12 -- $25\\micron$ emission) which requires an abundance of ultrasmall grains in diffuse regions. \\item The charge distribution of the small grains, which affects the rates of rotational excitation and damping. We have used standard estimates for photoelectric cross sections (Bakes \\& Tielens 1994) and electron capture cross sections (Draine \\& Sutin 1987) for very small grains. The uncertainties in these quantities affect both the electric dipole moment and, more importantly, the rate of angular momentum exchange with ions. \\item The electric dipole moments of both neutral and charged small grains. \\end{enumerate} The shape of the small grains (chainlike vs. sheetlike vs. quasispherical) is also uncertain, but is less critical than the above factors. A fifth factor -- the efficiency of $\\HH$ formation on small grains -- is not critical for these estimates: in Fig.\\ \\ref{fig:omega} only the grains in the CNM show appreciable sensitivity to whether or not $\\HH$ formation takes place, and even for the CNM component the emissivity at $\\nu \\gtsim 2\\GHz$ is only slightly changed (see Fig.\\ \\ref{fig:gamma_effect}). Because of these uncertainties, definitive predictions for the rotational emission spectrum are not yet possible. Nevertheless, within the existing theoretical and observational uncertainties it appears that much or all of the observed 15--90 GHz ``anomalous'' emission is due to spinning dust grains. The largest discrepancy between observation and theory is at 14.5 GHz, where Leitch et al. report emission about 3.5 times stronger than the emission predicted in Fig.\\ \\ref{fig:j_nu}. Additional measurements at $\\nu\\ltsim 20\\GHz$ will be of great value to clarify whether this emission has another origin, or whether some of our assumptions concerning the dust must be modified. Emission from rotating grains must be allowed for in studies of the cosmic microwave background radiation (see DL98). It appears that the small rotating grains may be partially aligned with the local magnetic field (Lazarian \\& Draine 1998), so that the electric dipole radiation will be linearly polarized; this may present a problem for interpretation of CMB polarization measurements by the MAP mission. } The principal results of this paper are as follows: \\begin{enumerate} \\item Even neutral dust grains are expected to usually have electric dipole moments arising from polarized chemical bonds within the grain. Charged grains have an additional contribution to the dipole moment due to displacement of the charge centroid from the mass centroid. Our estimate for the dipole moment is given by eq.\\ (\\ref{eq:mu2}). \\item The excitation and damping of rotation in small grains is determined by collisions with ions and neutrals, ``plasma drag'', emission of infrared and microwave radiation, and formation of H$_2$ on the grain surface. Ion collisions and plasma drag, omitted in previous estimates of rotation rates, are included in the present analysis and found to often dominate rotational excitation and damping. Induced-dipole attraction of neutrals by charged grains, and of ions by neutral grains, can also be significant. Because the charge state of the grain, and the fractional ionization of the gas, do not reflect thermodynamic equilibrium, the fluctuation-dissipation theorem does not directly apply. \\item For very small grains ($a\\ltsim7\\times10^{-8}\\cm$), the angular momentum of colliding ions is large compared to the r.m.s. angular momentum of the grain, and therefore the grain rotation history consists of ``rotational spikes'' separated by intervals of gradual rotational damping. \\item The estimated grain rotation rates are such that the small grains which have been postulated to explain the near-infrared emission feature can account for the emission observed at 30 -- 50 GHz by Kogut et al., de Oliveira-Costa et al., and Leitch et al. (1997). \\item The emission observed at 14.5 GHz by Leitch et al. is stronger than we estimate for spinning grains by a factor $\\sim4$. Additional determinations of emission from dust at frequencies $\\ltsim 20\\GHz$ will be of great value. \\item We predict that dark clouds should produce detectable microwave emission, with antenna temperatures of $\\gtsim 1 {\\rm mK}$ at $\\sim 10-30\\GHz$. \\end{enumerate}" + }, + "9802/astro-ph9802149_arXiv.txt": { + "abstract": "In the framework of the Kompaneets approximation the propagation of a shock front (SF) in inhomogeneous medium with the power-law density decrease at the exponent $n=2$ is investigated. For this important case corresponding to outer regions of the solar and stellar coronas an unexpectedly simple exact solution clearing a structure of the general solution for an arbitrary monotonic density of medium is found. Our results for plane-layered medium are compared to the Korycansky one for off-centre explosion in radially stratified medium. The relation between the solutions in these media are found and a new exact solution for a noncentral explosion in the case of density singularity on the finite radius is received. ", + "introduction": " ", + "conclusions": "" + }, + "9802/astro-ph9802180_arXiv.txt": { + "abstract": "We present a radio survey of X-ray sources in the Large and Small Magellanic clouds with the Australia Telescope Compact Array at 6.3 and 3.5 cm. Specifically, we have observed the fields of five LMC and two SMC supersoft X-ray sources, the X-ray binaries LMC X-1, X-2, X-3 \\& X-4, the X-ray transient Nova SMC 1992, and the soft gamma-ray repeater SGR 0525-66. None of the targets are detected as point sources at their catalogued positions. In particular, the proposed supersoft jet source RXJ 0513-69 is not detected, placing constraints on its radio luminosity compared to Galactic jet sources. Limits on emission from the black hole candidate systems LMC X-1 and X-3 are consistent with the radio behaviour of persistent Galactic black hole X-ray binaries, and a previous possible radio detection of LMC X-1 is found to almost certainly be due to nearby field sources. The SNR N49 in the field of SGR 0525-66 is mapped at higher resolution than previously, but there is still no evidence for any enhanced emission or disruption of the SNR at the location of the X-ray source. ", + "introduction": "Radio synchrotron and X-ray emission, though at opposite ends of the electromagnetic spectrum, are both tracers of high-energy phenomena in astrophysical sources. Thermal X-ray emission clearly demonstrates the existence of material at extremely high temperatures, while radio synchrotron emission originates in the spiralling of highly relativistic electrons around magnetic field lines. It has become increasingly apparent that the behaviour of sources in one energy regime may be correlated with that in another, although the exact mechanism is often unclear. In particular, radio and X-ray emission from compact Galactic X-ray sources are often related. While radio emission from such sources, typically at distances between 1 - 10 kpc, is now relatively routinely detected and monitored, it has yet to be detected from a binary in an extragalactic system. With this in mind, we have searched for radio emission from some of the most powerful X-ray emitters in the nearest external galaxies, the Large and Small Magellanic clouds, at $\\sim 55$ and $\\sim 60$ kpc respectively. Targets for our survey included most of the Magellanic cloud {\\em supersoft X-ray sources}, several bright LMC X-ray binaries and other transient systems. \\subsection{Supersoft sources} The prototypical supersoft X-ray sources, CAL~83 and CAL~87, were first detected in the Large Magellanic Cloud in 1979-1980 with the {\\it Einstein} X-ray Observatory (Long, Helfand \\& Grabelsky, 1981), although later {\\it ROSAT} observations have considerably enlarged the group (Tr\\\"{u}mper et al.\\ 1991). The defining characteristics of the supersoft sources are their extremely low X-ray energies and high bolometric luminosities (typically $L_{\\rm bol} \\sim 10^{38}$~erg\\,s$^{-1}$ and T$_{\\rm bb} \\sim$~tens of eV, where T$_{\\rm bb}$ is the blackbody temperature). Progress in determining the exact nature of these systems has been hindered by the fact that they are undetectable in the Galactic plane, due to the high level of soft X-ray absorption (e.g. van den Heuvel et al.\\ 1992). Most currently known systems are therefore optically faint extragalactic objects, predominantly in the Magellanic clouds and M31 (see e.g. Kahabka \\& Tr\\\"umper 1996 for a review). Although the term ``supersoft source'' has been previously applied to a range of objects such as planetary nebula nuclei (Wang 1991) and PG~1159 stars (Cowley et al.\\ 1995), we shall consider here only those objects exhibiting the characteristics of X-ray binaries (e.g. Crampton et al.\\ 1987; Smale et al.\\ 1988; Pakull et al.\\ 1988; Cowley et al.\\ 1990; Pakull et al.\\ 1993). The model for the SSSs which has gained predominanceis that of an accreting white dwarf in a binary system which is undergoing steady nuclear burning on its surface (van den Heuvel et al.\\ 1992). However, it should be noted that models of black hole (Cowley et al.\\ 1990; Crampton et al.\\ 1996) and neutron star accretors (Greiner, Hasinger \\& Kahabka 1991; Kylafis \\& Xilouris 1993) also exist. Indeed, one of the systems considered in this paper, RXJ~0059-71, almost certainly contains a neutron star, although, with 2.7~s pulsations (Hughes 1994) and a Be-type secondary star (Southwell \\& Charles 1996), it is not strictly an archetypal source. Of the other supersoft objects considered here, RXJ~0513-69 is unique in being the only SSS to exhibit optical jets (Pakull 1994, private communication; Cowley et al.\\ 1996). The source is an X-ray transient which was discovered in outburst during the {\\it ROSAT} All Sky Survey (Schaeidt, Hasinger \\& Tr\\\"{u}mper 1993). The optical spectrum (Pakull et al.\\ 1993; Cowley et al.\\ 1993; Crampton et al.\\ 1996; Southwell et al.\\ 1996) is similar to that of CAL~83, and the two sources have comparable optical magnitudes (V $\\sim 16-17$~mag). However, only RXJ~0513-69 exhibits Doppler-shifted components of He{\\sc ii}~4686 and H$\\beta$, with velocities characteristic of the escape speed of a white dwarf (Southwell et al.\\ 1996), implying the presence of a highly-collimated outflow. However, despite the drawing of analogies with SS433, neither this source nor any other supersoft X-ray binary have ever been detected at radio wavelengths. One of the LMC sources considered here, RX~J0550-71, does not yet have an optical counterpart, hence it should be noted that the nature of this object is particularly uncertain. \\begin{table*} \\begin{minipage}{120mm} \\caption{Radio survey of Magellanic Cloud X-ray sources with the Australia Telescope compact array. All upper limits are 3$\\sigma$.} \\begin{tabular}{cccccc} Source & Object &\\multicolumn{4}{c}{Point-source flux density (mJy)} \\\\ & type & 3.5 cm & 6.3 cm & 12.7 cm & 21.7 cm\\\\ \\hline RXJ 0513-69 & LMC supersoft source & $<0.09$ & $<0.09$ & $<0.15$ & $<0.18$ \\\\ RXJ 0528-69 & LMC supersoft source & $<0.09$ & $<0.09$ & &\\\\ CAL 83 & LNC supersoft source & $<0.12$ & $<0.12$ & &\\\\ CAL 87 & LMC supersoft source & $<0.12$ & $<0.12$ & & \\\\ RXJ 0550-71 & LMC supersoft source & $<0.15$ & $<0.15$ & & \\\\ LMC X-1 & BHC X-ray binary & $<1.5$ & $<1.5$ & & \\\\ LMC X-2 & X-ray binary & $<0.15$ & $<0.15$ & &\\\\ LMC X-3 & BHC X-ray binary & $<0.12$ & $<0.18$ & & \\\\ LMC X-4 & X-ray pulsar X-ray binary & $<0.15$ & $<0.18$ & &\\\\ SGR 0525-66 & Soft $\\gamma$-ray repeater & $<0.3$ & $<0.6$ & &\\\\ \\hline 1E 0035-72 & SMC supersoft source & $<0.12$ & $<0.12$ & & \\\\ RXJ 0059-71 & SMC supersoft source & $<0.12$ & $<0.15$ & & \\\\ Nova SMC 1992 & SMC X-ray transient & $<0.18$ & $<0.15$ & &\\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*} \\subsection{Radio emission from X-ray binaries} Radio emission has been detected from approximately 20\\% of Galactic X-ray binary systems, comprising a neutron star or black hole accreting matter from a more normal companion (e.g. Hjellming \\& Han 1995). In several cases the emission has been resolved by high-resolution observations into jet-like structures reminiscent of outflows from AGN, and relativistic or near-relativistic velocities inferred (e.g. Fender, Bell~Burnell \\& Waltman 1997 and references therein). It now seems that black hole systems are particularly likely sources of radio emission, whether transient or persistent -- in either case the characteristics of the radio emission are mirrored in those of the X-ray emission. Persistent black hole systems in quiescence appear to have centimetric radio luminosities which agree within a factor two with each other. For example, observations of Cygnus X-1, GX 339-4, 1E1740.7-2942 and GRS 1758-258 are all consistent with a centimetric flux density of $\\sim 10$ mJy at 3 kpc. More transient systems such as Cygnus X-3, GRS 1915+105 and X-ray `novae' such as V404 Cyg, show much more dramatic variability, with measured flux densities from $<1$ to $>20 000$ mJy. LMC X-1 and LMC X-3, two very luminous black hole candidates may be included in the class of persistent X-ray sources. At a distance of 55 kpc, we would only expect to observe a flux density of a few tens of $\\mu$Jy, by analogy with their Galactic cousins. However, lack of knowledge of their true nature and the chance of catching a rare flaring state makes the observations worthwhile. In the case of LMC X-1 this is particularly so, as Spencer et al. (1997) report the detection of a significant ($\\sim 80$ mJy) flux from the field of this source. LMC X-4 is an X-ray pulsar system, containing a highly magnetised accreting neutron star. Such systems in our Galaxy are found not to be radio-emitters (Fender et al. 1997). LMC X-2 is a low-mass X-ray binary system thought to contain a neutron star. While we considered both these sources far less likely to be detected than LMC X-1 or LMC X-3, they were included in the survey for completeness. Nova SMC 1992 (Clark, Remillard \\& Woo 1996) was discovered in archival ROSAT observations of 1992 Oct 1-2 as an extremely bright transient X-ray source. Clark et al. (1996) proposed a nearby 14th magnitude blue star as the optical counterpart, and suggested the source may be the first high-mass black hole X-ray nova detected. While previous X-ray `novae' have often been very bright radio sources (see above), they are also generally associated with low-mass companion stars and the nature of this system remains uncertain. \\subsection{SGR 0525-66} SGR 0525-66 is one of a group of only three confirmed (Smith 1997) and one possible additional (Hurley et al. 1997) soft $\\gamma$-ray repeaters. These are sources of repeated bursts of low-energy $\\gamma$-rays, possibly associated with neutron stars. SGR 0525-66 is coincident in the sky with the SNR N49, and a soft X-ray counterpart has been found (Rothschild, Kulkarni \\& Lingenfelter 1994). No radio, optical or infrared counterpart has been identified however (Dickel et al. 1995), although Rothschild et al. (1994) predicted a compact plerion-like radio nebula to be associated with the source. Any physical link with N49 has yet to be established. ", + "conclusions": "We have surveyed the fields of eight LMC and SMC supersoft X-ray sources, the X-ray binaries LMC X-1, X-2, X-3 \\& X-4, and the soft $\\gamma$-ray repeater SR 0525-66 at radio wavelengths. We have found no point-source radio emission from any of the sources. In particular we find no detectable radio emission from the proposed jet source RXJ 0513.9-6951. Neither do we detect nebulosity, such as that observed optically around CAL83, associated with any of the supersoft sources. Limits on emission from the black hole candidate X-ray binaries LMC X-1 and LMC X-3 are consistent with the radio brightnesses of their Galactic analogues. We show that a possible previous radio detection of LMC X-1 was almost certainly due to nearby field sources, and that due to its location in a radio-bright part of the LMC, this source is going to be very difficult to ever detect. Limits on radio emission from the other two X-ray binaries are as expected. The SNR N49, which SGR 0525-66 appears to lie on the northern edge of, shows no enhanced emission at the location of the SGR, nor any disruption to its structure suggesting association between the two. We can constrain the radio luminosity of any compact (arcsec-scale) structure associated with the SGR to be more than an order of magnitude below that which we might expect if the X-ray source discussed in Rothschild et al. (1994) were indeed a synchrotron nebula powered by the SGR. In summary, we have placed limits on radio emission from a variety of X-ray sources in the Magellanic Clouds, finding nore of them to be anomalously bright by comparison with the Galactic counterparts. The detection of radio emission from extragalactic X-ray binaries and related systems is likely to require an increase in the sensitivity of ground-based arrays or coordinated and fortunate observations during an outburst." + }, + "9802/hep-ph9802295_arXiv.txt": { + "abstract": "s{This talk summarizes our recent work which studied the impact of resonant neutrino conversion induced by some non-standard neutrino properties beyond mass and mixing, such as neutrino magnetic moment, lepton-flavor non-universality as well as flavor changing neutral current interactions in SUSY models with broken $R$ parity, on supernova physics. } ", + "introduction": "Neutrino flavor conversion could cause some significant influence on supernova physics \\cite{review}. In this talk we discuss the effect of such conversion induced by non-standard properties of neutrinos, not just by mass and mixing, on supernova physics. In particular, we consider the effect on % neutrino shock-reheating, supernova heavy elements nucleosynthesis as well as $\\bar{\\nu}_e$ signal, and show that in some case rather stringent limits on model parameters can be obtained. \\subsection{Some basic features of supernova neutrinos} A type-II supernova occurs when a massive star ($M \\gsim 8M_\\odot$) has reached the last stage of its life \\cite{suzuki}. Almost all of the gravitational binding energy of the final neutron star (about $\\sim 10^{53}$ erg) is radiated away in form of neutrinos. The individual neutrino luminosities in supernovae are approximately the same but the individual neutrino energy distributions are very different because they interact differently with the star material, as following reactions show, \\begin{eqnarray} \\label{nu-n} \\nu_e+n&\\to & p+e^-,\\\\ \\label{nu-p} \\bar\\nu_e+p& \\to & n +e^+,\\\\ \\label{nu-N} \\nu + N & \\to &\\nu + N, \\ \\ (N=p,n). \\end{eqnarray} Since the cross sections of the charged-current reaction is larger than that of the neutral-current one and there are more neutrons than protons, the $\\nu_e$'s have the largest interaction rates with the matter and hence thermally decouple at the lowest temperature. On the other hand, $\\nu_{\\tau(\\mu)}$ and $\\bar\\nu_{\\tau(\\mu)}$'s lack the the charged-current absorption reactions on the free nucleons inside the neutron star and hence thermally decouple at the highest temperature. As a result, the average neutrino energies satisfy the following hierarchy: \\begin{equation} \\label{hierarchy} \\langle E_{\\nu_e}\\rangle <\\langle E_{\\bar\\nu_e}\\rangle < \\langle E_{\\nu_{\\tau(\\mu)}}\\rangle \\approx\\langle E_{\\bar\\nu_{\\tau(\\mu)}}\\rangle. \\end{equation} Typically, the average supernova neutrino energies are, $\\langle E_{\\nu_e}\\rangle \\approx 11\\ \\mbox{MeV},\\ \\langle E_{\\bar\\nu_e}\\rangle \\approx 16\\ \\mbox{MeV},\\ \\langle E_{\\nu_{\\tau(\\mu)}}\\rangle \\approx \\langle E_{\\bar\\nu_{\\tau(\\mu)}}\\rangle\\approx 25\\ \\mbox{MeV}$. \\subsection{Impact of neutrino oscillation on supernova physics} Here we very briefly review some significant effects of neutrino oscillation, which occurs if neutrinos are massive and mixed, % on supernova physics, studied in some previous work. First issue is concerned with the neutrino conversion effect on shock re-heating in the delayed explosion scenario \\cite{delayed}. If neutrinos are massive and mixed and follow the mass hierarchy as observed in the quark sector, one expects MSW resonant conversion \\cite{MSW} between $\\nu_e$ and $\\nu_\\mu$ or $\\nu_\\tau$ inside supernova. If the conversion occurs between the neutrinosphere and the stalled shock this can help the explosion \\cite{fuller}. Due to the conversion the energy spectra of $\\nu_e$ and $\\nu_\\mu$ or $\\nu_\\tau$ can be swapped and hence $\\nu_e$ would have larger average energy leading to a larger energy deposition by reactions in eqs. (\\ref{nu-n}) and (\\ref{nu-p}) so that the stalled shock would be re-energized. Second issue is the impact on heavy elements nucleosynthesis in supernova. To have successful $r$-process the site must be neutron rich, i.e. $Y_e < 0.5$ where $Y_e$ is number of electron per baryon. The $Y_e$ value is mainly determined by the competition between the two absorption reactions in eqs. (\\ref{nu-n}) and (\\ref{nu-p}). In the standard supernova model the latter process is favoured due to the higher average energy of $\\bar\\nu_e$ which guarantees the neutron richness. If the neutrino oscillations do occur between the neutrinosphere and the region relevant for $r$-process the site can be driven to proton-rich due to the reaction (\\ref{nu-n}) and therefore, $r$-process could be prevented \\cite{qian}. Third argument is the effect of neutrino oscillation on $\\bar{\\nu}_e$ signal in the terrestrial detector. It is discussed that from SN1987A data \\cite{kamimb} the large oscillation between $\\bar{\\nu}_e$ and $\\bar{\\nu}_\\mu$ or $\\bar{\\nu}_\\tau$ is disfavoured since the oscillation can induce harder $\\bar{\\nu}_e$ spectra than the observed one \\cite{ssb}. Finally, although we will not discuss this further in this talk, we also mention that neutrino oscillation in the presence of polarization of the medium in the star due to the strong magnetic field could lead to some interesting consequences \\cite{pol}. ", + "conclusions": "We have discussed the impact of the resonant conversion induced by some non-standard properties of neutrinos. In particular we focussed on the conversion induced by neutrino transition magnetic moment, flavor non-universality, FCNC interaction in SUSY models with broken $R$ parity, and mixing with some sterile state. We have shown that some significant effects on supernova physics are expected and in some case we can derive bounds on neutrino parameters from the shock re-heating, $r$-process as well as SN1987A $\\bar{\\nu}_e$ signal arguments. We note that these bounds are first of all supernova model dependent, but complementary to the ones we obtain in the laboratory experiments and sometimes they are happen to be more stringent." + }, + "9802/astro-ph9802305_arXiv.txt": { + "abstract": "Recent observations have shown that the \\feka\\ line profile of the Seyfert~1 galaxy MCG-06-30-15 is strongly variable. We attempt accretion disk model fits to the \\feka\\ line profiles in a high, low and medium continuum luminosity phase of this source. During the monitoring by Iwasawa et al. (1996) a broad red-shifted component remained reasonably constant while a narrower component at $\\approx$ 6.4 keV strongly responded to continuum changes. Physically consistent fits are possible if the index $\\xi$\\ of the power-law emissivity changes from 0.7 (high phase) to 3.0 (low phase). The shape of the red-shifted component at low phase is crucial to the disk model interpretation. We suggest that the actual shape may be a broad redshifted Gaussian. Three lines of evidence support the interpretation of the \\feka\\ line as multicomponent, beyond the lack of correlation in the response to continuum changes of the red and blue components in MCG -06-30-15. (1) We show that the strong concentration of narrow peak centroids at 6.4 keV is inconsistent with expectations of a random distribution of disk orientations. (2) The average \\feka\\ profile for a sample of 16 mostly Seyfert~1's suggests a natural decomposition into two Gaussians one unshifted/narrow and the other redshifted/broad. (3) Evidence for emission in excess to the expectation of disk models on the high energy side of the \\feka\\ profile is both a challenge for low inclination disk models and support for the two component decomposition. ", + "introduction": "In an earlier paper (Sulentic et al 1997; hereafter S97) we considered the problems associated with models that see the X-ray \\feka\\ line emission at $\\approx$ 6.4 keV in Seyfert~1 galaxies as arising from fluorescence reflection (or emission) from an accretion disk. A broad and redshifted \\feka\\ emission feature is observed in the spectra of many Seyfert~1 galaxies. The situation for other AGN classes is less well defined and will not be considered here. Seyfert~1 emission profiles with a narrow unshifted peak at rest energy $\\approx$ 6.4 keV, and a broad red-shifted wing extending down to $\\approx$ 4.5 keV are most suggestive of an accretion disk line profile (see e.g. Tanaka et al. 1995). A great deal of effort has been expended towards observing them and fitting their spectra with disk models (see S97 for references). In S97 we attempted to build a disk illumination model that could simultaneously produce both \\feka\\ and optical Balmer lines. The motivation was to somehow constrain the wide dispersion in disk model fits that have been published for Seyfert~1 galaxies. While we could reproduce the profile widths and approximate emitting radii consistent with observations, our model fits to the observed Balmer line profiles were especially poor. In addition there were cases of severe disagreement between inclinations derived for \\hiha\\ and \\feka. The diversity in Seyfert~1 line profiles does not allow us to obtain a convergence in disk model parameter space. Nandra et al. (1997a,b) also found empirically that no single emissivity law could account for the diversity in line profiles among their sample of 16 mostly Seyfert~1's. These problems may be more easily overcome if one is prepared to consider a significant flux contribution from a second non--disk source. Recently a challenge to disk models has arisen from the most studied Seyfert~1 galaxy MCG-06-30-15 (Iwasawa et al. 1996; hereafter I96). The \\feka\\ line profile appears to change dramatically in response to continuum variations. In \\S \\ref{data} we summarize the variability data and its challenge to the line emitting accretion disk interpretation for \\feka\\ posed by MCG-06-30-15. We attempt a solution to all variability phases in \\S \\ref{expla}. In \\S \\ref{compo} we consider the evidence for the composite nature of the \\feka\\ profile, and in \\S \\ref{basta} we briefly discuss the implications for disk models. \\begin{figure} \\figurenum{1} \\plotone{fig01var.eps} \\caption[1]{ \\feka\\ line profiles (filled circles) observed by Iwasawa et al. and best fitting disk model profiles (filled line). See text and Table 1 for model parameters. Ascissa units are rest energy in keV; ordinate units are 10$^{-5}$ photons cm$^{-2}$ keV$^{-1}$ s$^{-1}$. Also shown (dashed line) is a gaussian profile.}\\end{figure} ", + "conclusions": "} Accretion disk model fitting to the variable \\feka\\ profile of MGC -06-30-15 require a Kerr black hole: even if only the low continuum phase \\feka\\ profile cannot be fitted by disk models around a Schwarzschild black hole, the angular momentum of a massive black hole obviously cannot change on timescales of $\\approx 10^5$ s. As the inner disk radius for a Kerr black hole with a/M = 0.998 is $\\approx$ 1.23 \\rg, and as the region where the ``Doppler boosted peak'' is formed occurs at $\\rm R \\approx 5-20$ \\rg, different illumination of the disk following changes in the continuum luminosity explains the strong change observed in the narrow peak at 6.4 keV (Plate 1 of Fanton et al. 1997 shows that it is not so for a Schwarzschild black hole). On the other hand, without attempting a physical interpretation of the \\feka\\ line profile, the strong variations of the narrow peak, along with the possibility of little or no variation in the redshifted broader line part, hint at two independent components. The two component hypothesis is reinforced by the detection of the narrow peak at 6.4 keV in the wide majority of cases. This statistical difficulty for accretion disk models must be understood, before accretion disk models can be accepted. The formulation of a physical model for a multicomponent \\feka\\ line goes beyond the aims of the present letter. We can, however, speculate on several possibilities. A possibility, recently revived by Misra \\&\\ Khembavi (1998), is that broadening could occur by Compton scattering the photons of an intrinsically narrower \\feka\\ line in a corona of size $\\approx$ 300 GM/c$^2$, much larger than the size of the region of continuum formation. Another possibility is that the narrow peak could be associated to the BLR: with an intrinsic width of 30000 \\kms\\, \\feka\\ emission from the BLR should appear as an unresolved peak. The maximum line shifts in the BLR are too small to be detected with ASCA, so that the peak would appear always at about 6.4 keV (as expected for cold iron emission). Clearly the detection of strong changes in the narrow peak with no apparent time delay (along with some difficulties raised by S97) challenges these interpretations, but a longer time coverage would be required to rule them out. The redshifted broader component may be, on the other hand, associated with the region of continuum production. Observations of a cut-off in the X-ray continuum of Seyfert galaxies at energy $\\approx$ 600 keV suggests the presence of a thermal corona above the surface of the accretion disk (e. g. Zdziarski et al. 1995; Haardt \\&\\ Maraschi 1991; also S97). \\feka\\ emission from highly ionized iron in the corona could be shifted to the observed rest energy by gravitational redshift." + }, + "9802/astro-ph9802133_arXiv.txt": { + "abstract": "We study the transition from radiation domination to matter domination in Jordan--Brans--Dicke theory, in particular examining how the Hubble length at equality depends on the coupling parameter $\\omega$. We consider the prospects for using high-accuracy microwave anisotropy and large-scale structure data to constrain $\\omega$ more strongly than by conventional solar system gravity experiments. ", + "introduction": "One of the most important epochs in the history of the Universe is the transition from radiation domination to matter domination. This transition alters the growth rate of density perturbations: during the radiation era perturbations well inside the horizon are nearly frozen but once matter domination commences, perturbations on all length scales are able to grow by gravitational instability. Consequently, the horizon scale at the time of matter--radiation equality is imprinted upon the spectrum of density perturbations; indeed, in a flat cold dark matter cosmology it is the only length scale to appear in the perturbation spectrum. The nature of this transition has been well studied in the standard general relativistic cosmology, and plays a crucial role in calculations of the perturbation spectrum and associated microwave background anisotropies by codes such as {\\sc cmbfast}~\\cite{SZ}. The success of general relativity as a description of our Universe allows us to evaluate the performance of rival theories of gravitation. So far, no weak-field or cosmological observations disagree with the predictions of general relativity. The most important class of deviant theories are scalar-tensor gravity theories, of which the Jordan--Brans--Dicke (JBD) theory \\cite{JBD,Will} is the simplest and best-studied generalization of general relativity. This theory leads to variations in the Newtonian gravitation ``constant'' $G$, and introduces a new coupling constant $\\omega$, with general relativity recovered in the limit $1/\\omega \\rightarrow 0$. The most robust constraint on $\\omega$, that it must exceed 500, has been derived from timing experiments using the Viking space probe \\cite{Reas}, and has stood for nearly 20 years now. Other constraints, such as those from nucleosynthesis \\cite{nucl}, are comparable but more model dependent; the most detailed analysis \\cite{nucl2} gives only $\\omega > 50$. Within the next five to ten years, the advent of new microwave anisotropy satellites, MAP and Planck, and large galaxy-redshift surveys, 2df and the Sloan Digital Sky Survey (SDSS), promises to revolutionize our understanding of cosmology by permitting the accurate determination of a large number of cosmological parameters \\cite{parest,Teg,parest2}. So far, estimates of the accuracy of parameter estimation have only been made for cosmological parameters, such as the Hubble constant $H_0$ and the matter density $\\Omega_{{\\rm m}}$, and for parameters describing the initial perturbations, such as the spectral index $n$. However, such techniques can also in principle be extended to constrain parameters defining the underlying gravity theory, such as $\\omega$, which also influence the gravitational instability process. A full analysis of the viability of obtaining such constraints is an imposing task; all the perturbation formalism employed to compute the present-day matter and radiation power spectra must be generalized to the theory of gravity under consideration, and the results then processed through the Fisher information matrix technology of Refs.~\\cite{parest,Teg,parest2}. In this paper, we assess whether such constraints might be competitive with existing bounds. We do this by studying the properties of the radiation--matter transition in JBD theory. We find that new cosmological data sets may well give limits competitive with those obtained from weak-field solar system tests of general relativity. ", + "conclusions": "We have studied the matter--radiation transition in the JBD theory, both numerically and analytically. The shift in the epoch of matter--radiation equality will influence the shape of the density perturbation spectrum, and it appears that precision microwave anisotropy measurements and large galaxy-redshift surveys may in the future be able to impose limits on $\\omega$ competitive with existing solar system bounds. However, it seems unlikely that a very substantial improvement will be possible. It may therefore be best to wait to see whether the high quality of data promised is actually delivered before embarking on the substantial undertaking of generalizing general-relativistic results to carry out a proper estimate of the likely observational limits. If all goes well, use of the data to constrain parameters of the gravitational theory will be a worthwhile endeavour and an unexpected bonus from future high-precision observational studies of galaxies and the microwave background." + }, + "9802/astro-ph9802243_arXiv.txt": { + "abstract": "Weak lensing effects are known to introduce non-linear couplings in the CMB temperature maps. In inflationary scenario, the primary CMB anisotropies are expected to form a 2D Gaussian map, for which, the probability distribution function of the ellipticity defined from the local temperature curvature matrix has a very specific shape. I show that lenses alter significantly the shape of this PDF, inducing an excess of elongated structures. The precise functional form is computed for both the field points and the temperature extrema. These analytical results are confirmed by numerical experiments on 10x10 square degree maps. These numerical results allow to investigate the effects of smoothing and to estimate the cosmic variance. For the best resolution and sky coverage of the Planck mission the signal to noise ratio for the statistical indicators presented here is about 3 to 6 depending on the cosmological models. A marginal detection should therefore be possible. ", + "introduction": "The detection by the COBE/DMR experiment (Smoot et al. 1993) of the anisotropies of the CMB temperature at very large angular scales, above $7\\deg$, has open an exiting new mean of investigation for cosmology. Many experiments that are now under development will provide us in the near future with precious data at smaller angular scale, down to a fraction of a degree. At such angular scales the dominant mechanism that generates the CMB anisotropies is the primary, linear order, coupling of the various cosmic fluids (photons, baryons and different species of possible dark matter) before and during recombinaison. Whether the primary anisotropies originate from quantum fluctuations in an inflationary scenario, or from other mechanisms such as topological defects generated in a phase transition epoch is yet unclear. This is undoubtlessly a major scientific goal for the coming experiments. Indeed, one of the clearest signature of the topological models is that they are expected to induce non-Gaussian primary temperature fluctuations (Pen et al. 1994, Turok 1996, Barnes \\& Turok 1996). But even in the case of inflationary scenario, it is possible that secondary effects, non-linear couplings of the radiation field with matter, induce non-Gaussian features. This is particularly important at very small angular scale, down to the arcmin scale. I am more particularly interested here in the static lens effects. The effects of the gravitational distortion on the power spectrum of the primary temperature maps have been the subject of many investigations in the last decade (Blanchard \\& Schneider 1987, Kashlinsky 1988, Cole \\& Efstathiou 1989, Sasaki 1989, Tomita \\& Watanabe 1989, Linder 1990, Cay\\'on, Mart\\'\\i nez-Gonz\\'alez \\& Sanz 1993a, b, Fukushige, Makino \\& Ebisuzaki 1994, Seljak 1996). It is now clear that the lens effect on the power spectrum is small, although it is probably worth to take it into account in a very detailed analysis of such measurements. In principle, for an inflationary scenario, a detailed analysis of the CMB power spectrum should allow to constrain very accurately the cosmological parameters (see for instance Jungman et al. 1996). But it would be anyway extremely interesting to be able to have a positive detection of the lens effect since it can potentially be used to constrain the amplitude of the cosmic density power spectrum, independently of the constraints obtained from the CMB power spectrum itself. It could therefore be a precious test for the cosmological model(s) favored by the shape of the CMB power spectrum. In a previous paper, the four point functions induced by the lens effects has been investigated (Bernardeau 1997). This is the most direct quantity that can be calculated from the coupling terms introduced by the lenses. One thus obtains specific properties of CMB maps induced by the lenses. Another possible way of investigation is to look for cross-correlation between the CMB temperature gradients and the displacement field induced by the projected large-scale structures (Suginohara et al. 1997). The signal however depends on the bias properties of the galaxies. It cannot be directly interpreted in terms of cosmological parameters. If the four point function is indeed sensitive to the intrinsic depth of the lens potential wells, it is not necessarily the best indicator in terms of signal to noise ratio for a realistic experiment. The aim of this paper is therefore to investigate new means of detecting these effects. In section 2 I present numerical results showing peculiar examples of the lens effects on temperature maps. The visual impression is that the lens effects induce a change in the topological properties of the temperature maps, and not so much in the local temperature distribution function. This is the motivation for the investigation of other statistical indicators that can reveal the lens effects with a better efficiency. Particularly interesting can be the statistical properties of the local curvature of the temperature maps. The matrix of the second order temperature derivatives has indeed specific statistical properties in case of a 2D Gaussian field. These properties are affected when the lens effects are taken into account because of the induced mode couplings. More specifically a quantity that can be a good tracer of the lens effect is the probability distribution function of the local ellipticity since lenses tend to systematically stretch the local temperature patches. Compared to the 4-point function the motivation for considering such a quantity is then double. First of all, the cosmic variance is expected to be smaller: It is given by statistical properties related to the second order derivative, and therefore more sensitive to the small scale temperature fluctuations. Secondly the effects of lenses on the four-point function were found to be proportional to the cosine of the angle joining the observational directions (see Bernardeau 1997), which tends to substantially reduce the lens signal when it is averaged for the computation of the four point moment. Such a cancellation is not expected for the distortion effect on the curvature. The derivation of the lens effect for the distribution function of the local ellipticity is presented in Section 3. In Section 4, results are confronted with numerical experiments. In particular I estimate the cosmic variance for the ellipticity statistics. In the last section I discuss the dependence of the lens effects on the cosmological parameters. ", + "conclusions": "\\subsection{The dependence on the cosmological parameters} \\begin{figure} \\vspace{6 cm} \\special{hscale=50 vscale=50 voffset=-185 hoffset=-30 psfile=OmegaLambda.ps} \\caption{The ratio $\\sigma_{\\kappa}^2(\\Omega,\\Lambda) \\Omega^{-1}/ \\sigma_{\\kappa}^2(\\Omega=1,\\Lambda=0)$ as a function of $\\Omega$ and $\\Lambda$ for a power law spectrum with $n=-1.5$.} \\end{figure} The dependence on the cosmological parameters is explored here in a quite rough way. The temperature power spectrum, and its relation to the density power spectrum, has indeed a rather complicated dependence on all the cosmological parameters. So in this simple study I will assume that the shapes of $C_l$ and $P(k)$ remain the same and simply discuss the dependence of the amplitude of $\\sigma_{\\kappa}$ on the cosmological parameters due to the density-convergence relationship. The ratio $\\sigma_{\\kappa}/\\sigma_8$ is obviously a quantity that one might want to consider but it should be compared to other constraints coming from large-scale structure formation. The number density of clusters can provide us with a constraint which is in principle free of bias contamination. It constrains the amplitude of the density fluctuations at roughly the $8\\,h^{-1}$ Mpc scale for an Einstein-de Sitter Universe (Oukbir et al 1997), which corresponds roughly to the scale of interest for the lens CMB effects. For open universes the dependence cannot be given simply in terms of $\\sigma_8$ because linear mass scale of the galaxy clusters is shifted to a larger scale. A dependence arises then with the slope of the power spectrum (Oukbir \\& Blanchard 1997). Ignoring these subtleties, and following Eke, Cole \\& Frenk (1996) I will simply assume that the actual constraint coming from the observed number density of rich clusters can be written, \\be \\sigma_8\\ \\Omega^{0.5}\\approx 0.6\\pm 0.1. \\ee It means that the larger $\\Omega$ the smaller $\\sigma_8$. To estimate the variation of $\\sigma_{\\kappa}$ with roughly a fixed number density of clusters, one should then compute $\\sigma_{\\kappa}^2(\\Omega,\\Lambda)\\ \\Omega^{-1}/ \\sigma_{\\kappa}^2(\\Omega\\!=\\!1,\\Lambda\\!=\\!0)$. The result is given in Fig. 5. One can see, as expected, that the magnitude of the signal is very weakly $\\Omega$ dependent if $\\Lambda=0$ but significantly grows with $\\Lambda$. It implies that the signal to noise ratio would be a factor about 2 larger for a $\\Lambda$-CDM. \\subsection{On the interest of lens effect detection.} The observational quantities that have been investigated here have been designed to be sensitive to the non-linear couplings induced by lens effects. It is of course possible to consider other quantities, such as the Minkowski functionals, whose general properties for a Gaussian field have been recently investigated in detail (Winitzki \\& Kosowski 1997 and Schmalzing \\& G\\'orski 1997). The lens effects might indeed significantly affect the Minkowski functional for 2D maps. For instance the shapes of high threshold peaks are shown to be more elongated, which implies that the averaged circumference near the top of the peak should increase. However, a complete theoretical investigation of this effect is quite difficult because this is a nonlocal indicator. And in general, the most efficient non-Gaussian indicators probably depend on the processes one wants to detect. The number of hot and cold spots seems to be a good way to detect topological defects. Is is obviously not the case for the lenses. In a previous paper the four-point function was considered. Actually the local curvature can be viewed as an other way to have access to the same information. Here the four point function is simply averaged in a way that avoid too much cancellation. And other possible quantity is the collapsed four-point function, that identifies with the four order cumulant of the local temperature. This quantity however tends to cancel the contribution of the various terms. This seems not to be the case for the local curvature. It does not mean however that we cannot do better for detecting the lens effects. For instance the two-point correlation function of the local ellipticity (with orientation) might be a good indicator. However one should have in mind that the number of CMB structures per lens is rather small. It is thus not obvious that it can improve the situation. Note that the results obtained for the lens effects, temperature four-point function or ellipticity statistics, are all sensitive to the lens two-point correlation function. It comes from the fact that the only quantity related to the lenses which is potentially available from CMB maps is its power-spectrum. The non-Gaussian properties of the lens population are for instance not accessible. If it is actually possible to detect a lens effect, the smallness of the signal to noise ratio with which it can be determined indicates that it will be pointless to use this information as a mean to constrain the cosmological parameters. But, it would be extremely interesting to be able to check that the amount of lens effects computed from models favored by the CMB power spectrum are in agreement with its detection, or its non-detection. Note however that other secondary effects might as well induce non-Gaussian properties. The Rees-Sciama effect in particular is due to the evolving non-linear potentials and has a source term which is intrinsically non-Gaussian. A peculiar case of this effect is the moving lens effect (Birkinshaw \\& Gull 1983): potential wells that move perpendicularly to the line of sight induce temperature fluctuations. This effect is quadratic with the cosmic fields (proportional to the local velocity times the gradient of the potential) and is thus intrinsically non-Gaussian. These cases are probably worth investigating. I expect however that they give a smaller non-Gaussian signal since the static lens effect is the only mechanism that couples the primary anisotropies to the line of sight potentials." + }, + "9802/astro-ph9802075_arXiv.txt": { + "abstract": "New age determinations of the galactic disk globular clusters 47~Tuc, M71 and NGC6352 have been performed with our up-to-date $\\alpha$-enhanced stellar models. We find that all three clusters are about 9.2 Gyr old and therefore coeval with the oldest disk white dwarfs. Several arguments are presented which indicate that the initial helium content of the stars populating these clusters is close to the solar one. We also revisit a total of 28 halo clusters, for which we use an updated [Fe/H] scale. This new metallicity scale leads on average to an age reduction of around 0.8 Gyr relative to our previous results. We compare the predicted cluster distances, which result from our dating method, with the most recent distances based on HIPPARCOS parallaxes of local subdwarfs. We further demonstrate that for the most metal-rich clusters scaled-solar isochrones no longer can be used to replace $\\alpha$-enhanced ones at the same total metallicity. The implications of the presented age determinations are discussed in the context of the formation history of the Galaxy. ", + "introduction": "The galactic globular clusters (GC) constitute the fossil record of the Galaxy formation epoch; their ages therefore provide fundamental informations about the timescale of the formation process, and put strong constraints on the age of the universe. In two recent papers (Salaris, Degl'Innocenti \\& Weiss 1997; Salaris \\& Weiss 1997; Papers I \\& II) we have redetermined ages for a sample of GCs by means of new stellar models employing the latest improvements in stellar input physics, noticeably in equation of state and opacities. In Paper I we demonstrated that the age of three extremely metal-poor clusters is around 12 Gyr due to the model improvements. In Paper II the same models were applied to a sample of 25 halo clusters. We put special emphasis on the careful application of two distance and reddening-independent methods for determining absolute and relative ages within groups of clusters of comparable metallicity. Not only did we confirm the new, low value for the age of the oldest clusters, but also found a clear correlation between age and metallicity of the halo clusters. The lowest age in our sample amounted to only 6.5 Gyr for Ter~7, which, however, might not be one of the typical halo clusters (see the discussion in Paper II), for which the lowest age was 8 Gyr. We emphasize that our predicted distances were in agreement with the first HIPPARCOS-based distance determinations. Several aspects of our results confirmed those of other papers (Chaboyer \\& Kim 1995; Mazzitelli, D'Antona \\& Caloi 1995) or have been confirmed in the following (Sarajedini, Chaboyer \\& Demarque 1997). Since the lowest ages for typical halo clusters were comparable to those of disk white dwarfs (Salaris et al.\\ 1997), it is interesting to ask how {\\sl disk} GC (see, e.g., Armandroff 1989 for a study about the kinematical properties of this GC system) fit into the emerging picture of the history of Galaxy formation. To determine their ages we have used the same rigorous approach we have been using successfully for the halo clusters. We selected three {\\sl disk} clusters, namely 47~Tuc, M71 and NGC6352, with $V$ and $B$ photometry extending well below the turn-off (TO) point, and with almost the same metallicity ([Fe/H]=-0.70 for 47~Tuc and M71, [Fe/H]=-0.64 for NGC6352 according to Carretta \\& Gratton 1997). In the case of 47~Tuc and NGC6352 the absolute age is directly determined by means of the difference between the Zero Age Horizontal Branch (ZAHB) level and the TO, the so-called $\\triangle V$ age indicator (or {\\sl vertical method}), while for M71, whose photometry shows a poorly populated HB, the relative age with respect to 47~Tuc is determined by means of the colour difference between the TO and the base of the Red Giant Branch (RGB), the so-called $\\triangle (B-V)$ age indicator (or {\\sl horizontal method}). Due to the high metallicity of the {\\sl disk} clusters, additional stellar evolution calculations became necessary. The computational details as well as the age determination method will be presented briefly in Sect.~2. The application of our isochrones to the three clusters will then be discussed in detail in Sect.~3. After that, we will revisit our halo cluster sample, which we extended by another three objects. Due to new results about the metallicity scale of globular clusters (Carretta \\& Gratton 1997), a re-evaluation of their ages became necessary. The updated results are contained in Sect.~4. In the final section we will discuss all results. ", + "conclusions": "The original motivation for Paper~I had been to resolve the apparent ``age discrepancy'' between globular clusters and the expanding universe. As we have demonstrated in all our papers, the oldest globular clusters appear to be younger than 12 Gyr, possibly only as old as 11 Gyr. Note that the neglect of diffusion even leads to a slight overestimation of the ages by approximately 0.8 Gyr (Cassisi et al.\\ 1998) for the vertical method (but the relative ages derived by means of the horizontal method are unchanged). Such a cluster age is completely consistent with the general range of recent $H_0$-determinations (see Freedman 1997 for a review), which extends from below 50 to 80 km/s/Mpc. Salaris \\& Cassisi (1998) used theoretical stellar models calculated with the same stellar evolution program to predict the brightness of the tip of the RGB (their tip bolometric luminosities are in complete agreement with our models), which can be used as a primary distance indicator. Applying this to the galaxy NGC3379 in the Leo~I group, the distance to the Coma cluster (obtained from the relative distance Coma-Leo I as derived from different secondary distance indicators) then gives a Hubble constant $H_0=60\\pm11$ km/s/Mpc. Depending on the cosmological model, the age of the universe is marginally or completely consistent with our oldest clusters. Since an open universe with $\\Omega=0.3$ and $\\Lambda=0$ appears to be preferred (the probably most important single evidence coming from the statistics of giant arcs produced by cluster of galaxies; see Bartelmann et al.\\ 1998), $H_0=60$ implies an age of $\\approx$13 Gyr. Therefore distance indicator and cluster age result in a completely consistent picture. \\begin{figure} \\includegraphics[scale=0.40,draft=false]{fig5pap3.ps} \\caption[]{The relation between age and metallicity for all clusters dated in Papers I, II and the present one (Table~2). Filled circles mark generic halo clusters, open circles clusters probably associated with satellite galaxies or affected by tidal interaction with the Magellanic Clouds. Crosses are the three disk clusters. For Rup106 and Pal12 ages obtained under the assumptions of $\\alpha$-element enhanced or scaled-solar metals are connected by thin lines.} \\protect\\label{f:afe} \\end{figure} \\begin{figure} \\includegraphics[scale=0.40,draft=false]{fig6pap3.ps} \\caption[]{ As in Fig.~\\ref{f:afe}, but this time the relation between age and galactocentric distances is shown.} \\protect\\label{f:adi} \\end{figure} The second purpose of our age determinations is to establish a reliable age-metallicity relation for galactic globular clusters. From Fig.~\\ref{f:afe} the general scenario of Paper~II is confirmed and extended: it appears that the more metal-poor clusters formed between 11 and 12 Gyr ago within a timespan of less than 1 Gyr. After that halo clusters continued to be formed at higher metallicity for about 4 Gyr, with a tendency for more metal-rich clusters to form later (with the exception of NGC6366, but see the discussion in Paper II about this cluster). The halo cluster creation continued even after the disk had already been formed, as it is evident from the ages of the oldest disk white dwarfs (Salaris et al.\\ 1997; Leggett et al.\\ 1997), which are around 9 Gyr old. In the present paper, we found that the three disk clusters investigated are coeval with these disk white dwarfs. This is an important result, because previous age estimates (e.g.\\ Chaboyer \\& Kim) placed disk clusters at the same age as our oldest halo clusters and let them appear to be considerably older than the disk white dwarfs. In addition, for all metallicities there are clusters (e.g.\\ Rup106 \\& Arp2) that appear to have been created at a later time, probably due to tidal interactions with dwarf spheroidal galaxies (namely the Sagittarius dSph) or the Magellanic Clouds. Such events are consistent with the finding that cluster formation in the LMC apparently has been triggered by close encounters with the Galaxy (Girardi et al.\\ 1995; Fujimoto \\& Kumai 1997). With respect to the age-galactocentric distance relation (Fig.~\\ref{f:adi}), our conclusion of Paper~II is supported: within the innermost 10 kpc an age spread of only $\\approx 2.5$ Gyr exists, if one neglects the exceptional clusters NGC6652 and NGC6366; the mean being at $\\approx$11 Gyr. The three disk clusters fit into this range at the lower boundary. At larger distances, the age differences approach $\\approx 5$ Gyr, with the mean age being somewhat smaller. We agree with the conclusion of Sarajedini et al.\\ (1997) in that there appears to be a tendency that the innermost parts of the galactic halo formed within a timespan shorter than that for the outermost regions. However, we find the age spreads to differ only by a factor of two, and given the small number of clusters, it is not clear whether this is significant enough. In particular, the larger spread for the outer halo clusters again depends on a few clusters. One of these is Pal12, for which there are indications (Brown et al.\\ 1997) that it is not $\\alpha$-enhanced. If true, its age would be raised by $\\approx 1.0$ Gyr (Sect.~4.1), making both the intermediate metal-rich group and the outer halo clusters more homogeneous in age. Another cluster with uncertain $\\alpha$-enhancement is Rup106. The disk clusters are different from the halo clusters with respect to their composition. We have found strong evidence that they have a solar-like helium content, while there is no indication of this in the halo clusters of comparable metallicity. Recalling that their metallicity is only half the solar one, this implies a different chemical enrichment history for them than for the solar neighbourhood. We also showed that $\\alpha$-element enhancement has to be taken into account properly in all aspects of the theoretical calculations for total metallicities $Z > Z_\\odot/10$, in order to obtain reliable isochrones. If the $\\alpha$-elements are enhanced in such a way that the condition [C+N+O+Ne/ Mg+Si+S+Ca+Fe]$\\approx$0 is violated, a fact that is compatible with current available observations, and the isochrone colours have to be determined with high accuracy for all evolutionary stages, then the computation of $\\alpha$-enhanced isochrones is necessary even for the lowest metallicities. Finally, distances predicted from our age determinations and ZAHB models are in very good agreement with HIPPARCOS-based data, demonstrating once more the reliability of our results." + }, + "9802/astro-ph9802061_arXiv.txt": { + "abstract": "I review the axionic solution of the strong CP problem and current status of the cosmic axion search. ", + "introduction": "Quantum chromodynamics before 1975 considered the following Lagrangian \\begin{equation} {\\cal L}=-{1\\over 2g^2}{\\rm Tr}F_{\\mu\\nu}\\tilde F^{\\mu\\nu} +\\bar q(iD^\\mu\\gamma_\\mu-M)q. \\end{equation} where $M$ is the diagonal, $\\gamma_5$-free, real quark mass matrix. But after 1975, the following term is known to be prensent in general in a world without a massless quark, \\begin{equation} +{\\bar\\theta\\over 16\\pi^2}{\\rm Tr}F_{\\mu\\nu}\\tilde F^{\\mu\\nu} \\end{equation} Since this $\\bar\\theta$ term violates CP invariance, the upper bound of the neutron electric dipole moment puts a strong constraint on the magnitude of $\\bar\\theta$, $|\\bar\\theta|<10^{-9}$. The smallness of $\\bar\\theta$ has led to the strong CP problem, $\\lq\\lq$Why is $\\bar\\theta$ so small?\"~\\cite{strcp} We know that many small parameters in physics have led to new ideas, in most cases leading to new symmetries. For example, $M_W/M_P\\ll 1$ has led to supersymmetry, $m_{u,d}\\ll 1$ GeV has led to $SU(2)_L\\times SU(2)_R$ chiral symmetry, etc. For the strong CP problem, the nicest solution is the very light axion resulting from the Peccei-Quinn symmetry~\\cite{pq}. ", + "conclusions": "" + }, + "9802/astro-ph9802257_arXiv.txt": { + "abstract": "Known millisecond pulsars have periods longer than $1.558$ ms. Recycled in binary systems, neutron stars can attain very short spin periods. In this paper we investigate the expected properties of the millisecond pulsar distribution by simulating synthetic populations under different assumptions for the neutron star equation of state and decay of the magnetic field. We find evidence that a tail in the distribution of millisecond pulsars may exist at periods shorter than those observed. ", + "introduction": "The shortest period at which neutron stars (NSs) have been observed ($P_{\\rm min}=1.558$ ms) is suspiciously close to the detection capabilities of pulsar (PSR) surveys carried out so far. A new experiment at the Northern Cross at Medicina has started recently whose sensitivity in the range $0.6 \\to 1.5$ ms is significantly better than that of the previous surveys (D'Amico et al., 1998\\markcite{x}). The experiment is aimed at determining whether $P_{\\rm min}$ is indeed a physical limit or is of instrumental origin. The minimum period of a stable, gravitationally bound NS is thought to be within $\\sim 0.6$ ms (for the softest equation of state of dense matter) and $\\sim 1.3$ ms (for the stiffest one). These are the limits derived by Cook, Shapiro \\& Teukolsky (1994a\\markcite{cst94l}; CST hereafter) in the context of a recycling scenario where accretion occurs from the inner edge of a Keplerian disk onto a bare unmagnetized NS. Remarkably, submillisecond periods are attained (before the mass-shedding instability is encountered) even in those models having a static maximum mass close to $1.4\\msole$. Several authors (Phinney \\& Kulkarni 1994\\markcite{pk94}; Stergioulas \\& Friedman 1995\\markcite{sf95}; Burderi \\& D'Amico 1997\\markcite{bd97}) discussed the importance of detecting an ultra-rapid rotating pulsar as a way of discriminating between the proposed equations of state (EoSs). The formation of a very fast spinning pulsar, however, depends sensitively on the history of the NS in the binary and on the evolution of the magnetic field which controls the dynamics of mass and angular momentum transfer at the magnetospheric boundary. Whilst detailed evolutionary calculations by Urpin, Geppert \\& Konenkov (1997) have shown that NSs move into the millisecond region of the plane $\\mu-P$ (magnetic moment versus period) under rather ``ordinary'' conditions, the possibility of populating the submillisecond range is still unknown. Though, on theoretical grounds, bare NSs can attain periods $P<1$ ms (CST), Nature may not provide the conditions for spinning a significant number of NSs to ultra short periods. In this Letter we investigate the efficiency of the recycling process in producing NSs in the yet unexplored region of the $\\mu-P$ diagram below $1.558$ ms. To this end, we simulate the distribution of a NS population evolved assuming three models for the magnetic field decay and selecting two representative EoSs. ", + "conclusions": "The results of our simulations are summarized in Figure 1, where we give the fraction of synthesized objects for EoS A and L in the 4 selected regions of the $\\mu~-~P$ plane. This fraction is derived normalizing the sample to the total number of pulsars spinning with a period $P<10$ ms. As a guideline, the upper left number in each cross represents the relative percentage of objects having $\\mu>10^{25.6}$ G\\cmtre (nowadays the minimum value detected in the observed population of millisecond pulsars) and a period $P<1.558$ ms (the minimum value observed). The typical variance is about $3\\%$ for the upper two quadrants of the cross. We find that {\\it NSs with periods below $1.558$} ms {\\it are present in a statistically significant number}, both in KB and BLOB models. As regard to the sensitivity of the results on the adopted parameters, we did not encounter major statistical differences when varying the torque function $g.$ The percentage of ultra rapid spinning PSRs is strongly depressed only when the relic magnetic field is an order of magnitude greater than that calculated in Konar \\& Bhattacharya (1997). As an illustration of the results of our runs, Figure 2 plots $\\mu$ versus $P$ for KB (top), BLOB and SIF (for EoS A). The first two scenarios give quite similar distributions below $P< 0.01$ s, whilst in SIF model, the population cluster between $0.01 \\to 1.00$ s, leaving the region below $1.558$ ms devoid of objects. In the BLOB model, NSs with relatively high $\\mu$ and long period populate the plane and this is a consequence of equation (6) predicting a bottom field higher than that found in KB, for the same $\\dot {m}$. In summary we find that: (a) The efficiency of the recycling process in spinning the NSs to the ultrashort periods can reach values as high as $20\\%$. {\\it This is the first indication that a tail in the distribution of the millisecond NSs may exist, at periods shorter than those effectively searched so far}. Our results provide an estimate of the relative importance of such a tail with respect to the bulk of the known millisecond PSRs. (b) As regard to the issue on the EoS, only in the BLOB scenario the distributions of ultra rapidly spinning PSRs enable us to discriminate between the EoS A and L. In the KB model, they are almost indistinguishable. A more pronounced difference in the distributions is found when we limit $M_{accr}$ to 0.5$\\msole$ for EoS L ($4\\% \\pm 1\\%$ for the left upper quadrant of KB model to contrast with the estimate of $9\\%\\pm3\\%$ for $M_{accr}=1\\msole$). (c) In the millisecond range, the distribution of $\\rm {Log} \\mu$ privileges values $> 25.6$. This comes from the hypothesis that the present millisecond pulsars evolved from the Low Mass X-Ray Binary population, with typical observed accretion rates ${\\dot{m}}_{accr}$ clustering around $0.1.$ If future observations will be in favor of low values of $\\mu$, it will be suggestive of the presence of a yet undetected population of low luminosity low mass binaries. (d) Pulsar search experiments with a sensitivity profile almost flat in the ultrashort period range, like the new one in progress at the Northern Cross radiotelescope (D'Amico et al. 1997), are strongly encouraged, as they might help putting constraints on the equation of state and on the evolutionary paths of recycled neutron stars." + }, + "9802/astro-ph9802127_arXiv.txt": { + "abstract": "We present the results of a comprehensive theoretical investigation on the Period-Radius (PR) relation of classical Cepheids based on new sequences of full amplitude, nonlinear convective models constructed by adopting a wide range of both stellar masses and chemical compositions. In the period range $ 0.9 \\le log P \\le 1.8$ a very good agreement is found between theoretical predictions and current available data whereas outside this range, both at shorter and at longer periods, nonlinear radii attain intermediate values between empirical relations based on different Baade-Wesselink (BW) methods and photometric bandpasses. ", + "introduction": "The BW method (Baade 1926; Wesselink 1946) has been receiving growing attention from the astronomical community since it allows direct measurement of both radii and absolute magnitudes. Even though some physical assumptions of this method were questioned (Karp 1975; Gautschy 1987; Bono, Caputo \\& Stellingwerf 1994; Butler et al. 1996), in the last years a paramount effort has been undertaken for improving its accuracy and consistency (Barnes \\& Evans 1976; Sollazzo et al. 1981; Laney \\& Stobie 1995, hereinafter LS; Ripepi et al. 1997, hereinafter RBMR). At the same time, Krockenberger, Sasselov \\& Noyes (1997, hereinafter KSN) have recently developed a new BW method, based on Fourier coefficients, for evaluating the uncertainty on mean stellar radii due to individual measurement errors. Substantial improvements in the measurements of both Cepheid radii and distances were thoroughly discussed in several outstanding papers by LS and more recently by Laney (1997), Di Benedetto (1997) and Gieren, Fouqu\\`e \\& G\\`omez (1997, hereinafter GFG). Despite this ongoing observational effort, theoretical investigations devoted to the Cepheid PR relation based on up-to-date evolutionary and pulsational models are lagging. In fact LS, by comparing the PR relation derived for a sample of 49 Galactic Cepheids with Fernie's (1984) weighted mean theoretical PR relation, found that the slope of the empirical relation is steeper than the theoretical one and that BW radii are 12\\% smaller than the theoretical ones for a period equal to 10 d. On the other hand RBMR, by adopting a new version of the CORS method (Sollazzo et al. 1981, and references therein) found, as expected (LS), that the slope of their PR relation is slightly shallower if compared either with empirical BW relations based on IR photometry, or with theoretical relations. The reason why so far only few investigations have been devoted to the evaluation of the mean theoretical PR relation is that its slope depends on the intrinsic width of the instability strip. The cool edge of the instability strip can be evaluated only by coupling the local conservation equations with a nonlocal and time-dependent equation for turbulent-convective motions (Stellingwerf 1982; Gehmeyr 1992). Theoretical PR relations available in the literature (Karp 1975; Cogan 1978) are based on radiative models and therefore cannot be considered \"pure\" theoretical relations. In fact, radiative models can only fix the location of the blue edge, whereas the temperature width of the instability strip is inferred from observational data. As a consequence, both the zero-point and the slope of these \"semi-theoretical\" PR relations depend on the completeness of the adopted sample and on the relations used for transforming the mean colors into mean effective temperatures. Moreover, Karp's and Cogan's relations have been derived by assuming that the width of the instability strip is constant when moving from short to long-period Cepheids. However, this assumption is not supported by observational estimates, and indeed Pel (1980) in a seminal investigation showed that the Cepheid instability region is not a rectangular-shaped but a wedge-shaped strip, i.e. the color range narrows toward short-period Cepheids. The main aim of this investigation is to establish the Cepheid PR relation on a genuine theoretical basis by adopting the mean radii and the periods predicted by full amplitude, nonlinear convective models and then to compare theoretical with empirical PR relations. ", + "conclusions": "We developed a new theoretical scenario of the actual properties of classical Cepheids in the Galaxy and in the MCs. By adopting both radii and periods predicted by full amplitude, nonlinear, convective models we found that the use of two different ML relations based on canonical and noncanonical (mild overshooting) evolutionary models has a marginal effect on the PR relation, and indeed in the mean PR relation the difference is of the order of 3\\%. At the same time, we also found that an increase in the metal content implies a decrease in the mean radius. This effect is not constant but increases when moving from short to long-period Cepheids. In particular, a change in the chemical composition from Y=0.25, Z=0.004 to Y=0.28, Z=0.02 implies at $log P\\approx2$ a decrease in the mean radius of the order of 9\\%. This result prompts that, within the current accuracy of both photometric and spectroscopic data, the dependence of the PR relation on metallicity could be detected and measured if a proper number of long-period variables -$P>40$ d- are included in the sample. Theoretical and empirical radii are found in very good agreement in the period range $0.9 \\le log P \\le 1.8$, but present some discrepancies toward short and long-period Cepheids. No firm conclusion was reached on the intimate nature of this discrepancy since current mean stellar radii estimated by adopting different BW methods, photometric bands, and data sets present a large scatter both at $log P < 0.7$ and $log P > 1.8$. Comparison between theory and observations suggests that the value of the {\\em p}-factor could change when moving from short to long-period Cepheids. At the same time, the results of this investigation disclose a new approach for testing the internal accuracy and the consistency of the assumptions adopted by the different BW methods. In fact, observables predicted by nonlinear, convective models can be fed to the progeny of the BW method for assessing the intervening effects of systematic errors and/or of possible biases in the radius measurements. It is a pleasure to thank D. Laney as a referee for his clarifying comments and valuable suggestions on current observational data. \\pagebreak" + }, + "9802/astro-ph9802194_arXiv.txt": { + "abstract": "Radial velocity measurements have proven a powerful tool for finding planets in short period orbits around other stars. In this paper we develop an analytical expression relating the sensitivity to a periodic signal to the duration and accuracy of a given set of data. The effects of windowing are explored, and also the sensitivity to periods longer than the total length of observations. We show that current observations are not yet long or accurate enough to make unambiguous detection of planets with the same mass and period as Jupiter. However, if measurements are continued at the current best levels of accuracy (5 m/sec) for a decade, then planets of Jovian mass and brown dwarfs will either be detected or ruled out for orbits with periods less than $\\sim$15 years. As specific examples, we outline the performance of our technique on large amplitude and large eccentricity radial velocity signals recently discussed in the literature and we delineate the region explored by the measurements of 14 single stars made over a twelve year period by Walker et al. (1995). Had any of these stars shown motion like that caused by the exo-planets recently detected, it would have been easily detected. The data set interesting limits on the presence of brown dwarfs at orbital radii of 5--10 AU. The most significant features in the Walker et al. data are apparent long term velocity trends in 36 UMa and $\\beta$ Vir, consistent with super planets of mass of 2 $M_J$ in a 10 year period, or 20--30 $M_J$ in a 50 year period. If the data are free of long term systematic errors, the probability of just one of the 14 stars showing this signal by chance is about 15\\%. Finally, we suggest an observing strategy for future large radial velocity surveys which, if implemented, will allow coverage of the largest range of parameter space with the smallest amount of observing time per star. We suggest that about 10-15 measurements be made of each star in the first two years of the survey, then 2--3 measurements per year thereafter, provided no (or slow) variation is observed. More frequent observations would of course be indicated if such variations were present. ", + "introduction": "For nearly two decades most high precision radial velocity surveys of nearby stars were focused on detecting radial velocity variations in stars due to companions with mass and period of Jupiter. The signature would consist of changes in the relative stellar radial velocity with a period of a decade and amplitude of a few tens of meters per second or less, depending on orbital inclination with respect to the solar system. The surprising recent result, triggered by the discovery of 51~Peg~B by Mayor \\& Queloz (1995), has been the finding that as many as 5-10\\% of solar type stars have companions with mass $<10 M_J$ and with periods less than $\\sim$3 years. No sub-stellar companions with periods longer than $\\sim$3 years have so far been detected by radial velocity searches. Are Jupiter mass companions at longer periods rare, or is it simply the case that current observations do not have the length or sensitivity to see them? Is the theoretical prediction by Boss (1995) correct, that Jovian planets should form preferentially at $>$4-5 AU separations from their primary? Our purpose in this paper is to show what we can learn from velocity data of a given duration and accuracy, to help plan continued programs. The best measurement errors for a series of radial velocity measurements so far published are those of Butler et al. (1996), who observe a magnitude $V=5$ star and quote an accuracy of 3 m/s for measurements taken over one year. Measurements up until this work have been limited to a lower accuracy standard of about 15 m/s. Several other programs (see section \\ref{strategy} for a list of radial velocity search programs currently underway) are planning new or expanded searches with a goal of obtaining measurements with similar accuracy. In light of these efforts, and in expectation of their eventual success in obtaining such accuracy, we shall use 5 m/s as a `canonical' value for the error in many of the examples and the discussion below. Such advances in radial velocity calibration allow accuracy to be relatively free from systematic error. Poissonian photon noise remains as the fundamental limit to accuracy. In this limit, strong constraints can be placed upon the existence of periodic radial velocity signals in a given set of data, given a suitable analysis technique. Many efforts have been made to determine whether a given set of data contains a signal. Most of those in common use are based upon the periodogram analysis techniques discussed by Scargle (1982). This technique is shown to be equivalent to a least squares fit for the signal at a given period, and he derives an exponential probability distribution of obtaining a false alarm from a given set of data. Horne and Baliunas (1986 hence HB) have refined the technique by showing that this exponential must be normalized to the total variance of the data and derived an empirical expression for the number of independent frequencies available to a set of data. Further refinements (Irwin et al. 1989, Walker et al. 1995) account for variable weighting of individual data points and correlations between fitted parameters. Our work represents a different approach in which, rather than dealing with least squares minimization indirectly through a periodogram analysis, we examine the best fits to the data directly and determine their significance. We derive an analytic expression for the probability that a given best fit velocity amplitude is non-random. We first develop analytical expressions relating sensitivity to planetary companions of different masses and periods, given velocity measurements of specified accuracy, duration and number. Motion with periods longer than the duration of observations is detected with reduced sensitivity, and this reduction is explored by Monte Carlo methods. We illustrate our analysis technique by application to the published set of radial velocity data from Walker et~al. (1995), the longest time baseline survey so far published, with quoted precision of 15 m/s. Limiting our analysis to the subset of 14 stars which have no known visual binary companion, we obtain quantitative upper limits to companions masses for orbital periods of a few days to periods as long as 100 years. Finally, we suggest a strategy for efficiently implementing a search of a large number of stars for radial velocity signatures due to the presence of a companion. ", + "conclusions": "" + }, + "9802/astro-ph9802089_arXiv.txt": { + "abstract": "The dramatic discovery with the {\\em Rossi X-Ray Timing Explorer\\/} satellite of remarkably coherent $\\sim$300--1200~Hz oscillations in the X-ray brightness of some sixteen neutron stars in low-mass binary systems has spurred theoretical modeling of these oscillations and investigation of their implications for the neutron stars and accretion flows in these systems. High-frequency oscillations are observed both during thermonuclear X-ray bursts and during intervals of accretion-powered emission and appear to be a characteristic feature of disk-accreting neutron stars with weak magnetic fields. In this review we focus on the high-frequency quasi-periodic oscillations (QPOs) seen in the accretion-powered emission. We first summarize the key properties of these kilohertz QPOs and then describe briefly the models that have been proposed to explain them. The existing evidence strongly favors beat-frequency models. We mention several of the difficulties encountered in applying the magnetospheric beat-frequency model to the kilohertz QPOs. The most fully developed and successful model is the sonic-point beat-frequency model. We describe the work on this model in some detail. We then discuss observations that could help to distinguish between models. We conclude by noting some of the ways in which study of the kilohertz QPOs may advance our understanding of dense matter and strong gravitational fields. ", + "introduction": "It has long been expected (see, e.g., \\cite{ELP86,LP79}) that important information about the intrinsic properties of neutron stars and stellar mass black holes, as well as about the physics of accretion onto them, could be extracted from their X-ray variability at frequencies comparable to the \\hbox{$\\sim$1--10~kHz} dynamical frequencies near them. The {\\em Rossi X-ray Timing Explorer\\/} (\\rxte\\/) was specifically designed \\cite{BS89,Sw95} to have the large area, microsecond time resolution, and high telemetry rate needed to probe this high-frequency regime. \\begin{figure}[t!] % \\vglue0.2truecm \\centerline{\\hskip0.25truecm \\psfig{file=rome.fig1.eps,angle=0,width=6.6truecm}} \\vglue-1.0truecm \\caption{\\label{fig1} Side view of an atoll source (top) and a Z source (bottom). The dark shading indicates the accretion disk. The light shading indicates the hot gas in the magnetosphere and the central corona that surrounds the neutron stars. The arrows in the lower panel indicate the cooler, approximately radial inflow that is thought to be present outside the central coronae of the Z sources. From [28]. } \\label{nsLMXBs} \\end{figure} The value of access to these high frequencies was dramatically confirmed when, less than two months after the launch of \\rxte, remarkably coherent $\\sim 1000$~Hz oscillations were discovered (see \\cite{vdK97b}) in the X-ray brightness of two neutron stars in low-mass X-ray binaries (LMXBs). These are the fastest astrophysical oscillations ever discovered. It is thought that they are being generated at or near the surfaces of these neutron stars. To date, brightness oscillations with frequencies ranging from $\\sim$300~Hz to more than 1200~Hz have been discovered in some sixteen neutron stars in LMXBs. High-frequency oscillations are observed both during thermonuclear X-ray bursts \\cite{Stroh96c,Stroh97,Swank98} and during intervals of accretion-powered emission \\cite{vdK97b,vdK98}. They appear to be a characteristic feature of disk-accreting neutron stars with weak magnetic fields. In this review we focus on the high-frequency quasi-periodic oscillations (QPOs) seen in the accretion-powered emission. The weak-field neutron stars in which these {\\em kilohertz QPOs\\/} have been discovered were studied extensively with \\exosat\\ and \\ginga, and a detailed physical picture was developed based on their 2--20~keV X-ray spectra and 1--100~Hz X-ray variability (see \\cite{GL92,HK89,L89,L91}). In this picture, LMXBs with weak-field neutron stars fall into two classes, called the ``Z'' and ``atoll'' sources after the shapes of the paths they trace, over time, in X-ray color-color diagrams \\cite{HK89}. Modeling \\cite{PL97,PL98,PLM95} of the X-ray spectra of these sources indicates that both types are surrounded by a central corona with a scattering optical depth $\\sim$3--10, and that the Z sources also have a cooler radial inflow (see Fig.~\\ref{nsLMXBs}). The six known ``Z'' sources have mass accretion rates comparable to the Eddington critical rate $\\mdote$ and inferred surface magnetic fields \\about$10^{9}$--$10^{10}\\,$G. They display two different types of lower-frequency QPOs. The $\\sim$15--50~Hz ``horizontal-branch oscillations'' (HBOs) \\cite{vdK85} have frequencies that increase steeply with accretion rate and are thought to be produced by the magnetospheric beat-frequency mechanism \\cite{AS85,GL92,L85}. The $\\sim$4--8~Hz ``normal/flaring branch oscillations'' (N/FBOs) \\cite{MP86} are thought to be caused by radiation-hydrodynamic oscillations in the radial inflow \\cite{FLM89,L89,L91}. The $\\sim$15 known ``atoll'' sources are both less luminous and more weakly magnetic than the Z sources, with accretion rates $\\sim$1--10\\% of $\\mdote$ and magnetic fields $\\sim 10^7 \\dash\\, 5\\ee9\\,$G. No QPOs with frequencies $\\lta 100\\,$Hz have so far been detected in any of the atoll sources, with the exception of Cir~X-1, which has a $\\sim$1--30~Hz QPO with a frequency that varies with its brightness (see \\cite{Bradt98}). Considerable theoretical effort has been devoted to understanding the mechanisms that produce the kilohertz QPOs. In this review we first summarize the key properties of these QPOs and then describe briefly the models that have been proposed to explain them. The kilohertz QPOs commonly occur in pairs, and the existing evidence strongly favors beat-frequency models. We discuss application of the magnetospheric beat-frequency model to the kilohertz QPOs. The most fully developed and successful model is the so-called sonic-point beat-frequency model, in which the higher frequency in a QPO pair is the orbital frequency of gas at the inner edge of the Keplerian disk flow and the lower frequency is the difference between this frequency and the spin frequency of the neutron star. We outline the sonic-point model and describe some of the calculations that have been carried out to explore it. We conclude by discussing observations that could help to distinguish between models and noting some of the ways in which the study of kilohertz QPOs may advance our understanding of neutron stars, dense matter, and strong gravitational fields. ", + "conclusions": "The discovery with \\rxte\\/ of $\\sim$1000~Hz brightness oscillations from a large number of accreting neutron stars is a spectacular achievement that validates both the scientific expectations that led to the mission and the long years of hard work that were needed to bring it to fruition. As explained in \\S3, the observations made to date strongly favor beat-frequency models of the kilohertz QPOs detected in the persistent emission. The magnetospheric beat-frequency interpretation of the kilohertz QPO pairs can account for the fact that their separation frequencies are approximately constant and related to the frequencies of the brightness oscillations seen during X-ray bursts, but this interpretation suffers from many serious difficulties. The sonic-point beat-frequency model also explains this fact and is the most fully developed and successful model, but it is not yet confirmed. Therefore, it is important to consider which further observations and calculations would be particularly helpful in testing the beat-frequency hypothesis and discriminating between these two beat-frequency models. The detection in the persistent emission of stable oscillations with frequencies harmonically related to those seen during X-ray bursts would confirm the spin-frequency interpretation of the latter and strongly support the beat-frequency interpretation of the kilohertz QPO pairs. On the other hand, detection of stable oscillations with frequencies not harmonically related to those seen during X-ray bursts would seriously undermine both models. Beat-frequency models also predict that weak oscillations will be present at other special frequencies, such as the first overtone of the beat frequency and the sum of the spin frequency and the orbital frequency, and detection of any of these would provide strong support for these models. It is therefore important to search for QPOs at these frequencies. Even if none are detected, upper limits on their amplitudes would be valuable, because they would constrain models of the coronae around these neutron stars (see Fig.~\\ref{nsLMXBs}). The sonic-point model predicts that the stronger the stellar magnetic field, the weaker the kilohertz QPOs will be \\cite{MLP97}. In the magnetospheric interpretation, one expects just the opposite. Figure~\\ref{qpoAmps} shows that the current data on kilohertz QPO amplitudes clearly favors the sonic-point model, but more precise and uniform amplitude measurements, as well as further progress in modeling the X-ray spectra of the Z and atoll sources, would help to clarify the situation. \\begin{figure}[t] % \\vglue-0.75truecm \\centerline{\\hglue0.75truecm\\psfig{file=rome.fig6.eps,height=8.7truecm,width=8.7truecm}} \\vglue-1.9truecm \\caption{ Comparison of $\\eta_{\\rm flow}$, the fraction of the specific angular momentum of an element of gas in Keplerian circular orbit at Boyer-Lindquist radial coordinate $r$ (measured in units of the stellar mass) that must be removed in order for the gas to fall from $r$ to the radius $R_{\\rm ms}$ of the innermost stable circular orbit, with $\\eta_{\\rm rad}({\\rm max})$ the largest fraction of the specific angular momentum of an element of gas that can be removed by radiation coming from the surface of a nonrotating, isotropically radiating, spherical star of radius $5M$. } \\label{AngMom} \\end{figure} Besides their intrinsic interest, the $\\sim$1000~Hz oscillations discovered with \\rxte\\ provide a new tool with which to probe the nature of strong gravitational fields and the properties of dense matter. Modeling of general relativistic effects on the gas dynamics and radiation transport processes involved in the generation of kilohertz QPOs may provide evidence for the existence of innermost stable circular orbits \\cite{MLP97,MLP98} and frame-dragging \\cite{ML93,ML96,Stella98,SV97}, both of which are important predictions of strong-field general relativity. If the frequencies of the highest-frequency kilohertz QPOs are orbital frequencies, as in beat-frequency models, these QPOs provide important new constraints on the masses and radii of the neutron stars in LMXBs and on the equation of state of neutron star matter \\cite{KFC97,MLP97,MLP98,ZSS97}. The brightness oscillations observed during X-ray bursts may constrain the compactness of neutron stars \\cite{ML98,Stroh97,Swank98}. New \\rxte\\ observations are continuing to yield significant fresh insights, and the rapid pace of important new discoveries is therefore likely to continue for many years. \\vskip6pt This work was supported in part by NSF grants AST~93-15133 and AST~96-18524 and NASA grant NAG~5-2925 at the University of Illinois, and NASA grant NAG~5-2868 at the University of Chicago." + }, + "9802/astro-ph9802040_arXiv.txt": { + "abstract": "Following a recent report that AO~Psc has broad iron \\kalpha\\ emission lines we have looked at the \\asca\\ spectra of 15 magnetic cataclysmic variables. We find that half of the systems have \\kalpha\\ lines broadened by \\sqig200 \\ev, while the remainder have narrow lines. We argue that the Doppler effect is insufficient to explain the finding and propose that the lines originate in accretion columns on the verge of optical thickness, where Compton scattering of resonantly-trapped line photons broadens the profile. We suggest that the broadening is a valuable diagnostic of conditions in the accretion column. ", + "introduction": "The accretion shock near the surface of a white dwarf in a magnetic cataclysmic variable (MCV) will consist of a highly ionized \\sqig10 \\kev\\ plasma cooling by bremsstrahlung emission. Most elements will be fully ionized, leaving iron as the dominant cause of line emission. While model fits to the X-ray spectra of MCVs have traditionally included an iron \\kalpha\\ line (e.g.~Norton, Watson \\&\\ King 1991), it is only with the \\asca\\ satellite that the data have sufficient spectral resolution to investigate the structure of the line. An analysis of the \\asca\\ data on AO~Psc (Hellier \\etal\\ 1996) showed that the thermal components of the iron \\kalpha\\ line were broadened by \\sqig150 \\ev. However, reports on two other MCVs with the same detectors [Kallman \\etal\\ (1996) on BY~Cam and Fujimoto \\&\\ Ishida (1997) on EX~Hya] found a \\kalpha\\ complex compatible with narrow emission lines. We have therefore investigated the line widths in the class as a whole. In this paper we present a systematic analysis of lines in 15 MCVs, and discuss the mechanisms responsible for line broadening. Of our sample, AM~Her is a phase-locked system (i.e.~a polar; see Cropper 1990 for a review), BY~Cam is nearly phase-locked, and the remaining 13 are non-phase-locked systems (i.e.~intermediate polars; reviewed by Patterson 1994). ", + "conclusions": "From the above discussion we conclude that the broad \\kalpha\\ iron lines found in roughly half of the MCVs studied are caused by a mixture of Doppler broadening due to radial infall, Compton down-shifted line emission from a reflected component, and, probably most importantly, Compton scattering of the line emission in the accretion column. A significant optical depth to scattering in the column may also be required to explain the line fluxes (e.g.~Swank \\etal\\ 1984; Done \\etal\\ 1995). We have suggested that the broadening originates in the transition region between optical thickness and optical thinness; that systems in which this transition occurs at a temperature too low for significant \\kalpha\\ emission have narrow lines; and that systems with broad lines must have regions of column which are still optically thick at a higher temperature (\\sqiggt 3 \\kev). We can test this by comparison with other work, taking the well studied stars EX~Hya, as an example of a system with narrow lines, and AO~Psc, the system with the clearest line broadening. In EX~Hya the presence of lines of elements with a lower ionization than iron implies a transition to optical thickness at $<$\\,1 \\kev\\ (Fujimoto \\&\\ Ishida 1997). In AO~Psc the line ratios imply a higher temperature transition [Fujimoto \\&\\ Ishida (1995) quote $<$ 3 \\kev, although this estimate is less certain due to the weaker lines] in agreement with the above reasoning. Further, we can detect optical thickness in the accretion column by looking at changes in the X-ray continuum as the white dwarf spins. From spin-resolved \\asca\\ spectroscopy of AO~Psc, Hellier \\etal\\ (1996) found that the column contained several phases of absorption. The densest, affecting regions of the column emitting at energies of at least 8 \\kev, requires an electron scattering column which changes by 6\\pten{23}\\ cm\\mintwo\\ over the spin cycle, suggesting an actual column of \\sqig2\\pten{24}\\ cm\\mintwo, and thus an optical depth of \\sqig 1. This column is compatible with an accretion rate of 10\\up{17}\\ g s\\minone\\ and an accretion area covering $10^{-3}$ of the white dwarf surface (Hellier \\etal\\ 1996), values which are in line with current estimates for intermediate polars (e.g.~Patterson 1994; Hellier 1997). In contrast, EX~Hya shows much less absorption, and in \\asca\\ data has no spin modulation above 6 \\kev\\ (Ishida, Mukai \\&\\ Osborne 1994; Allan, Hellier \\&\\ Beardmore 1998), implying that it is optically thin throughout the hard X-ray emitting regions. We can extend this difference into a test of our explanation, and predict that the other systems with clearly narrow lines, such as V1223~Sgr, will also be optically thin in the hard X-ray emitting regions. Thus spin-resolved spectroscopy with \\asca\\ would not show a counterpart of the very dense absorber revealed in AO~Psc. Unfortunately the test isn't clear cut since we need to distinguish between a flux reduction caused by an optically thick accretion column, and a flux reduction caused by emitting regions passing over the limb of the white dwarf. However, at any one pole, occultation effects and absorption effects are likely to occur in anti-phase in MCVs (e.g.~Hellier, Cropper \\&\\ Mason 1991) allowing the test to be performed given \\asca -quality spectroscopy and an understanding of the spin pulse in each system. Analysis of X-ray spectroscopy of other MCVs is thus needed to confirm these ideas." + }, + "9802/astro-ph9802330_arXiv.txt": { + "abstract": "}[2]{{\\footnotesize\\begin{center}ABSTRACT\\end{center} \\vspace{1mm}\\par#1\\par \\noindent {\\bf Key words:~~}{\\it #2}}} \\newcommand{\\FigCap}[1]{\\footnotesize\\par\\noindent Fig.\\ % \\refstepcounter{figure}\\thefigure. #1\\par} \\newcommand{\\TabCap}[2]{\\begin{center}\\parbox[t]{#1}{\\begin{center} \\small {\\spaceskip 2pt plus 1pt minus 1pt T a b l e} \\refstepcounter{table}\\thetable \\\\[2mm] \\footnotesize #2 \\end{center}}\\end{center}} \\newcommand{\\Table}[3]{\\begin{table}[htb]\\TabCap{#2}{#3} \\vspace{#1}\\end{table}} \\newcommand{\\TableSep}[2]{\\begin{table}[p]\\vspace{#1} \\TabCap{#2}\\end{table}} \\newcommand{\\TableFont}{\\footnotesize} \\newcommand{\\TableFontIt}{\\ttit} \\newcommand{\\SetTableFont}[1]{\\renewcommand{\\TableFont}{#1}} \\newcommand{\\MakeTable}[4]{\\begin{table}[htb]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}} \\newcommand{\\MakeFrameTable}[5]{\\begin{table}[htb]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1}\\hline\\trule #4\\\\[1mm] \\hline\\trule #5\\\\[1mm]\\hline \\end{tabular}\\end{center}\\end{table}} \\newcommand{\\MakeOwnTable}[6]{\\begin{table}[htb]\\TabCap{#2}{#3} \\begin{center}\\TableFont #4 \\begin{tabular}{#1} #5 \\end{tabular} #6 \\end{center}\\end{table}} \\newcommand{\\MakeTableSep}[4]{\\begin{table}[p]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}} \\newcommand{\\Figure}[2]{\\begin{figure}[htb]\\vspace{#1} \\FigCap{#2}\\end{figure}} \\newcommand{\\FigureSep}[2]{\\begin{figure}[p]\\vspace{#1} \\FigCap{#2}\\end{figure}} \\newcommand{\\FigureOwn}[3]{\\begin{figure}[htb]\\vspace{#1}#3 \\FigCap{#2}\\end{figure}} \\newcommand{\\FigureOwnSep}[3]{\\begin{figure}[p]\\vspace{#1}#3 \\FigCap{#2}\\end{figure}} \\newcommand{\\LeavePage}{\\begin{table}[p]\\vspace{19.5cm}\\end{table}} \\newenvironment{references}% { \\footnotesize \\frenchspacing \\renewcommand{\\thesection}{} \\renewcommand{\\in}{{\\rm in }} \\renewcommand{\\AA}{Astron.\\ Astrophys.} \\newcommand{\\AAS}{Astron.~Astrophys.~Suppl.~Ser.} \\newcommand{\\ApJ}{Astrophys.\\ J.} \\newcommand{\\ApJS}{Astrophys.\\ J.~Suppl.~Ser.} \\newcommand{\\ApJL}{Astrophys.\\ J.~Letters} \\newcommand{\\AJ}{Astron.\\ J.} \\newcommand{\\IBVS}{IBVS} \\newcommand{\\PASP}{P.A.S.P.} \\newcommand{\\Acta}{Acta Astron.} \\newcommand{\\MNRAS}{MNRAS} \\renewcommand{\\and}{{\\rm and }} {Results of the first two month of observations using the All Sky Automated Survey prototype camera are presented. More than 45 000 stars in 24 Selected Fields covering 140 sq. degrees were monitored a few times a night resulting in the $I$-band catalog containing $10^7$ individual measurements. Period search revealed 126 periodic variables brighter than 13 mag. Only 30 of them are known variables included in GCVS. The other 90 objects are newly detected variables - mainly eclipsing binaries (75\\%) and pulsating stars (17\\%). We estimate that completeness of the current catalogs of variable stars is smaller than 50 \\% already for the stars brighter than 9 mag. The Catalog is accessible over the WWW: {\\em http://www.astrouw.edu.pl/$\\sim$gp/asas/asas.html} }{Catalogs -- Stars:variables -- Surveys} \\vspace*{-6pt} ", + "introduction": "The All Sky Automated Survey (Pojma{\\'n}ski 1997, hereafter Paper I) is a new observing project which ultimate goal is detection and investigation of any kind of the photometric variability present all over the sky (Paczy{\\'n}ski 1997). We want to achieve this aim at relatively low cost, using simple automatic modules. In 1997 we have started monitoring over 20 Selected Fields to the limiting magnitude 13 (in $I$-band) using prototype automated mount equipped with 768 $\\times$ 512 MEADE Pictor~416 CCD camera, 135~mm f/1.8 telephoto lens and $I$-band (Schott RG-9, 3mm) filter. The instrument was placed at the Las Campanas Observatory which is operated by the Carnegie Institution of Washington, in the vicinity of the new OGLE-2 telescope (Udalski, Kubiak and Szyma{\\'n}ski 1997), where room for the control computer was kindly allocated. During the routine observations control program loops over the list of selected fields pointing camera and taking 3 minute exposures. Dark and flat-field images are exposed at the beginning of the night and appropriate data reduction process is applied after data acquisition in the fully automated way. Results of the aperture photometry are put into the ASAS Catalog, from which they may be retrieved e.g. using the World Wide Web. Detailed description of the prototype instrument, data acquisition and reduction process and ASAS Catalog can be found in Paper I. This paper presents results of the search for periodic variables in the ASAS Catalog using data obtained during the first two month of the prototype instrument operation. ", + "conclusions": "\\input table5.tex The Catalog of the Short Period Variable Stars in the Selected Fields contains 126 objects found among 45000 stars. About 90 of them where not previously known to be variable. Results presented in Table 5 suggest that completeness of the existing catalogs is smaller than 50\\% already for stars brighter than 9 mag and drops significantly for fainter objects. This drop is even more pronounced if one takes into account, that current ASAS survey starts to be incomplete at about 11~mag. There are many more other variable stars in the Selected Fields. Many of them are long period variables, which were already discovered, but their period has not been determined yet. We will perform detailed analysis of such objects using the data collected during almost one year of the ASAS prototype instrument operation. Our preliminary results show, that the small-scale instruments are ideal tools for reducing incompleteness of our knowledge about bright objects on the sky. We are going to extend our survey for much larger area of the sky, trying to increase the ASAS Catalog completeness to 13~mag using longer focal length. The Catlog of the Short Period Variable Stars, as well as the ASAS Catalog are accessible over the World Wide Web:\\\\ \\centerline{\\em http://www.astrouw.edu.pl/$\\sim$gp/asas/variables.html}\\\\ or\\\\ \\centerline{\\em http://www.astrouw.edu.pl/$\\sim$gp/asas/asas.html}. \\vspace*{-6pt} \\Acknow{It is a great pleasure to thank Prof.\\ Bohdan Paczy{\\'n}ski for initiative in this project, valuable discussions and providing necessary funds. We are indebted to the OGLE collaboration for letting us use facilities of the Warsaw telescope and for the permanent support of our instrumentation (opening and closing enclosure, tape exchanging and many more). Special thanks are due to Andrzej Udalski for his invaluable technical help and reanimating our equipment after serious break-down, to Micha{\\l} Szyma{\\'n}ski for lots of computing hints and to Marcin Kubiak and Przemek Wo{\\'z}niak for instrument re-adjustment after the earthquake. I am very indebted to Carnegie Institution of Washington for providing the excellent site for my instrument. This work was partly supported by the KBN BST grant.}" + }, + "9802/astro-ph9802106_arXiv.txt": { + "abstract": "The extragalactic microlensing scenario for natural wormholes is examined. It is shown that the main features of wormhole lensing events upon the light of distant Active Galactic Nuclei (AGNs) are similar to some types of already observed Gamma Ray Bursts (GRBs). Using recent satellite data on GRBs, an upper limit to the negative mass density -- ${\\cal O}\\,(10^{-36})$ g cm$^{-3}$ -- under the form of wormhole-like objects is presented.\\\\ {\\it PACS number(s): 98.62.Sb, 04.20.Gz} \\hspace{5.2cm} SUSSEX-AST-98/1-2 ", + "introduction": "Ten years after the seminal paper by Morris and Thorne \\cite{motho}, we face the following situation: there is no observational evidence supporting the existence of natural wormholes nor serious theoretical reasons for its impossibility \\cite{VISSER-BOOK}. Black holes shared such a status during years until the discovery of galactic X-ray sources and quasars in the 1960s. Wormholes, entities that warp spacetime in such way as to provide shortcuts to separated regions of the universe or even a way to allow a backward time travel, require the violation of the energy conditions (technically speaking, the null energy condition) in order to exist. The energy conditions are conjectures that are widely used to prove issues concerning singularities and black hole thermodynamics; they constitute just plausible statements, like the positivity of the energy density. However, several situations in which the energy conditions are violated are known; perhaps the most quoted of them is the Casimir effect. These violations are typically very small (of order $\\hbar$) and it is far from clear whether there could be macroscopic quantities of this kind of {\\it exotic} matter. Nevertheless, there is nothing really compelling to prevent its occurrence and wormholes might naturally exist \\cite{6-visser,7-visser}. Very recently, the consequences of the energy conditions were confronted with possible values of the Hubble parameter and the gravitational redshifts of the oldest stars in the galactic halo \\cite{VISSER-HUBBLE}. It was deduced that for the currently favored values of $H_0$, the strong energy condition should have been violated sometime between the formation of the oldest stars and the present epoch. On the other hand, negative gravitational masses (underdensities in the primordial universe) have been proposed as an explanation of the voids observed in the extragalactic space \\cite{PIRAN-0}. An early universe cosmic network of wormholes has also been suggested as an alternative solution for the cosmological horizon problem \\cite{Hochberg}. Mann \\cite{MANN} have found, in addition, that dense regions of negative mass can undergo gravitational collapse, ending up in exotic black holes that could populate the universe contributing to the bulk of total dark matter. All these works clearly show that it is at least possible that natural wormholes or other negative mass objects might exist. Then the study of their possible observational effects deserves serious consideration. Although no universal mechanism to generate a relic density of exotic matter is well established at present (because of our ignorance of quantum gravity laws), several interesting ideas have been recently proposed in the literature, like, for instance, the enlargement of \\wh throats --via inflation-- from the quantum foam to macroscopic sizes \\cite{ROMAN}. Despite current theoretical speculations suggest that the existence of compact objects of negative mass is plausible, their amount has not been yet constrained by observations. To provide such a constraint is the main goal of this paper. As far as we are aware, the first observational proposal to search for natural wormholes or similar gravitational negative anomalous compact objects was presented by Cramer et al. \\cite{CRAMER} (see also Ref. \\cite{DIAZ}). They suggested that gravitational microlensing effects of these objects upon the light of background stars could produce MACHO\\footnote{MACHO: massive compact halo object.}-like events \\cite{PACHINSKY-ANNU-REVIEW}, although with different (asymmetric) temporal profiles. Partial analysis of the results of several ongoing microlensing monitoring programs seems to show that wormhole-like objects are not present in the dark halo of our galaxy. (Hereafter, when speaking of negative masses, we shall think in this ingredient as always threading a wormhole. Although this can be relaxed for the development and analysis of the ideas to be considered, we shall do it just because it can provide useful numerical estimates and a pretty theoretical framework). In this paper we shall study the microlensing scenario for an extragalactic natural wormhole acting upon light coming from an Active Galactic Nucleus (AGN). It will be shown that such anomalous microlensing event would produce lightcurves very similar to some already observed Gamma Ray Bursts (GRBs) \\cite{BATSE-ANNU-REVIEW} and that this can be used to constrain the amount of negative mass in the universe. Preliminary results on this issue were introduced in \\cite{CONCURSO} and briefly commented on in \\cite{NS}. The paper is ordered as follows. The next section will review the relevant observational characteristics of the GRB phenomenon. Sec. III will deal with the negative mass lensing formalism. Afterwards, we shall analyze the consequences of negative-mass microlensing with an AGN as background source in Sec. IV. The possible nature of the lenses is treated in Sec. V, while the BATSE database is briefly discussed in the Sec. VI. The final two sections deal with the cosmological consequences of a negative mass distribution of compact objects and the conclusions. ", + "conclusions": "We have shown that microlensing events produced by wormholes with AGNs as background sources very much resemble certain types of GRBs; types that cannot be explained in standard models. We then used observational data on GRBs to determine an upper limit for the amount of wormhole-like objects in the universe. This upper limit is enough to see that negative matter hardly would have any influence in cosmology. An unusual feature of the presented scenario is that, while GRB repetition has previously been seen as a strong evidence for noncosmological origin, the microlensing model accepts it warmly: sources are cosmological and repetitions arise from different caustic crossings. This model implies that not only some bursting events must repeat, but also that they should do it with temporal profiles of specular character. This makes the model capable to be falsified. We expect that, with the improvement of the observational techniques and the increase of the GRB sample, more exact limits to the amount of the negative mass will be available. Forthcoming technologies and satellites such as the {\\it High Energy Transient Explorer} (HETE), the next {\\it Gamma Ray Large Area Space Telescope} (GLAST) and the current Beppo-SAX satellite will help to improve burst position measurements yielding light onto the repetition phenomenon. Whether the laws of physics, in some deep realization, forbid the violations of the energy conditions in the large amount needed to produce stellar-size compact objects of negative matter, is something not yet clear. But, if the universe does admit wormholes geometries in it, it is very likely that some of the GRBs may be caused by a microlensing mechanism, being this one of the main conclusions of this work. As an immediate spinoff, we have the converse fact, i.e. that if there were no possible burst in a large, perhaps not already obtained sample, which could be associated with wormhole-like lensing, then it should be understood as a serious objection to the existence of anomalous compact objects in the universe." + }, + "9802/nucl-th9802024_arXiv.txt": { + "abstract": "A semi--microscopic model for the low--energy photodisintegration of the $^9$Be nucleus is constructed, and the experimental data are analyzed with its help. The older radioactive isotope data are supported by this analysis. The theoretical photodisintegration cross section is derived. The astrophysical rates for the reaction $ \\alpha+\\alpha+n\\rightarrow ^9$Be$+\\gamma$ and the reverse photodisintegration of $^9$Be are calculated. The new reaction rate for $\\alpha+\\alpha+n\\rightarrow ^9$Be$+\\gamma$ is compared with previous estimations. ", + "introduction": "Recently fully microscopic calculations of nuclei with $A\\le 9$ have become feasible \\cite{pudl97,varga96}. The $^9$Be nucleus is such a system of special interest, as it allows tests of theories of interaction of composite particles \\cite{varga96}. The properties of low--energy continuum of $^9$Be are of particular importance in this connection. However, the corresponding experimental data on the low--energy photodisintegration of $^9$Be are not in mutual agreement (see Fig.~1). In the present work we develop a semi--microscopic model to describe the process, and we analyze the experimental data with its help. The model accounts simultaneously for both resonant and non--resonant contributions to the cross section. An estimation of the reliability of various data is obtained and a theoretical photodisintegration cross section is derived. We also calculate the reaction rates of the reaction $^9$Be$+\\gamma\\rightarrow\\alpha+\\alpha+n $ and the reverse reaction for astrophysical conditions. These reaction rates are of relevance in the high--entropy bubble in type II supernovae, an astrophysical site that has been suggested for the r--process \\cite{woo94,tak94}. The baryonic matter in this bubble is dominated in the beginning by $\\alpha$--particles, neutrons, and protons. The abundance distribution shifts then to higher masses through the recombination of the free $\\alpha$--particles, neutrons, and protons. This generates the so--called $\\alpha$--process leading to the formation of massive isotopes ($A \\approx 100$). The reaction path in the $\\alpha$--process is mainly determined by requirements of nuclear statistical equilibrium and depends also on the reaction rates of the various recombination paths bridging the mass 5 and 8 gaps. It has been shown that there are three principal reaction paths from $^4$He to $^{12}$C:\\\\ (i) $^4$He(2$\\alpha$,$\\gamma$)$^{12}$C\\\\ (ii) $^4$He($\\alpha$ n,$\\gamma$)$^9$Be($\\alpha$,n)$^{12}$C\\\\ (iii) $^4$He(2n,$\\gamma$)$^6$He($\\alpha$,n)$^9$Be($\\alpha$,n)$^{12}$C. It was shown in Refs.~\\cite{tak94,woo92,gor95} that the triple--alpha process (i) can be neglected compared to the reaction sequence (ii) via $^9$Be under r--process conditions in the $\\alpha$--process. Also the reaction path (iii) via $^6$He can be neglected for this scenario \\cite{gor95,efr96}. This is true even if the reaction rate of $^4$He(2n,$\\gamma$)$^6$He is strongly enhanced~\\cite{her97}, because then $^6$He is also destroyed very effectively through photodissociation. Therefore, for the $\\alpha$-- and r--process the reaction $^4$He($\\alpha$ n,$\\gamma$)$^9$Be plays a key role in bridging the unstable mass gaps at $A=5$ and $A=8$. The reaction rates of $^4$He($\\alpha$ n,$\\gamma$)$^9$Be and the reverse photodisintegration of $^9$Be were estimated in Ref.~\\cite{fowler75} from the experimental photodisintegration cross section. However, Ref.~\\cite{fowler75} did not include information on which experimental data their estimate was based. In view of the astrophysical relevance of these reactions we recalculate in the present work the rates of the first step of the reaction (ii) above. The same problem is also addressed in Ref.~\\cite{gor95}. These authors obtain the resonant contribution to the $^9$Be$(\\gamma,n)^8$Be cross section from the Breit--Wigner formula for the first excited state of $^9$Be with the parameters taken from Ref.~\\cite{ajz88}.\\footnote{We note that the $\\Gamma$ and $E_0$ parameters of the resonance used in Ref.~\\cite{gor95} seem to be incorrect. The resonant properties of the $1/2^+$ state of $^9$Be will be considered in our future work.} In order to calculate the non--resonant contribution they introduce a single--particle potential with the depth chosen to reproduce the ground state, calculate both ground-- and final--state continuum wave functions in this potential, and multiply the cross section obtained by the shell--model spectroscopic factor. They then add this cross section constructively or destructively to the resonant cross section to establish possible upper and lower bounds for the reaction rates. This procedure has certain shortcomings: a resonant contribution to the cross section should not emerge as an addition to the dynamic model used, since a correct quantum mechanical model should necessarily contain such a contribution itself, along with the non--resonant contribution and an interference term. Besides, the potential wells used for the ground state and continuum state should in fact be different: an additional spin--orbit potential, for example, should be present in the ground p--state as compared to the continuum s--state. In our model we use a three-body specification of the $^9$Be bound state, and a semimicroscopic continuum wave function which describes the essential scattering degrees of freedom at low relative energies. In Sect.~2 this model is formulated. In Sect.~3 the results for the $^9$Be$(\\gamma,n)$ cross section are given. In Sect.~4 the astrophysical rates for $\\alpha+\\alpha+n\\rightarrow ^9$Be$+\\gamma$ and the reverse reaction are calculated. ", + "conclusions": "" + }, + "9802/nucl-th9802030_arXiv.txt": { + "abstract": "\\noindent The equation of state and the properties of neutron stars are studied for a phase transition to a charged kaon condensate. We study the mixed phase by using Gibbs condition with comparison to the hitherto applied Maxwell construction. Implications for kaon condensation and for the mass-radius relation of condensed neutron stars are examined. ", + "introduction": " ", + "conclusions": "" + }, + "9802/astro-ph9802218_arXiv.txt": { + "abstract": "We examine the effects of cooling flows on the $T_{\\rm X}-L_{\\rm Bol}$ relation for a sample of the most X-ray luminous ($L_{\\rm Bol} > 10^{45}$ \\ergps) clusters of galaxies known. Using high-quality ASCA X-ray spectra and ROSAT images we explicitly account for the effects of cooling flows on the X-ray properties of the clusters and show that this reduces the previously-noted dispersion in the $T_{\\rm X}-L_{\\rm Bol}$ relationship. More importantly, the slope of the relationship is flattened from $L_{\\rm Bol} \\propto T_{\\rm X}^3$ to approximately $L_{\\rm Bol} \\propto T_{\\rm X}^2$, in agreement with recent theoretical models which include the effects of shocks and preheating on the X-ray gas. We find no evidence for evolution in the $T_{\\rm X}-L_{\\rm Bol}$ relation within $z \\sim 0.3$. Our results demonstrate that the effects of cooling flows must be accounted for before cosmological parameters can be determined from X-ray observations of clusters. The results presented here should provide a reliable basis for modelling the $T_{\\rm X}-L_{\\rm Bol}$ relation at high X-ray luminosities. ", + "introduction": "The baryonic content of clusters of galaxies is dominated by the metal-rich, X-ray luminous intracluster medium (ICM) that pervades the cluster potentials. X-ray observations permit precise measurements of the X-ray luminosities and temperatures ($T_{\\rm X}$) of clusters, which can be related to the masses of these systems. For self-similar clusters with a characteristic density that scales as the mean density of the universe, $T_{\\rm X} \\propto M^{2/3}(1+z_f)$, where $z_f$ is the redshift of formation of the cluster (e.g. Kaiser 1991; Evrard \\& Henry 1991). Since the bolometric luminosity is dominated by bremsstrahlung emission at X-ray wavelengths, we can then show that $L_{\\rm Bol} \\propto f^2 T_{\\rm X}^2 (1+z_f)^{3/2}$, where $f$ is the mass fraction in X-ray gas. (The redshift dependency holds for an Einstein-de-Sitter universe and is modified for lower values of $\\Omega_0$; \\eg Eke, Navarro \\& Frenk 1997). In principle, studies of the $L_{\\rm Bol}-T_{\\rm X}$ relation may be used to examine the evolution of clusters and should, for example, enable the range of cluster formation redshifts to be estimated (\\eg Scharf \\& Mushotzky 1997). It is well known that the observed $L_{\\rm Bol}-T_{\\rm X}$ relation for clusters contains a significant intrinsic dispersion. Within the simple theoretical framework outlined above, this dispersion should relate to the ages of the clusters. However, such simple models do not account for important physical processes that could effect the X-ray properties of clusters and, in particular, the effects of radiative cooling. In the central regions of most ($70-90$ per cent) clusters, the cooling time of the ICM is significantly less than the Hubble time (Edge \\etal 1992; Peres \\etal 1997), which leads to the formation of cooling flows (Fabian 1994; Allen \\& Fabian 1997). Fabian \\etal (1994) showed that clusters with cooling flows are offset in the ($2-10$ keV) $L_{\\rm X}-T_{\\rm X}$ plane such that they have lower temperatures for a given luminosity (or alternatively higher luminosities for a given temperature). Before concluding that cooling-flow clusters are simply older than non-cooling flow systems, however, one must account for the effects of the cooling gas on the X-ray properties of the clusters. Cooling flows can account for up to 70 per cent of the total X-ray luminosity of a cluster (Peres \\etal 1997; Allen \\etal 1998) and contain gas with a wide range of densities and temperatures. A spectral fit to a cooling-flow cluster with a simple isothermal model will yield only a mean emission-weighted temperature that will be reduced with respect to the virial value. Evidence for this effect was presented by Allen, Fabian \\& Kneib (1996), from a study of the nearby, luminous cooling-flow cluster PKS0745-191. These authors showed that explicitly accounting for the effects of the large cooling flow in this cluster leads to an X-ray determined mass value in excellent agreement with gravitational lensing measurements, but as much as a factor 3 greater than the mass inferred from a simple single-temperature X-ray analysis. Allen (1998) presented results for a larger sample of clusters and showed that for all of the cooling flow clusters in that sample the virial temperatures, determined from multiphase X-ray analyses which accounted for the effects of the cooling flows, were significantly higher than those inferred from simple single-temperature models and lead to consistent X-ray and lensing mass measurements. A second issue relating to the $L_{\\rm Bol}-T_{\\rm X}$ results is the slope of the relation, which appears to fit $L_{\\rm Bol}\\propto T_{\\rm X}^3$ (\\eg Edge \\& Stewart 1991; David \\etal 1993; Fabian \\etal 1994) rather than $L_{\\rm Bol}\\propto T_{\\rm X}^2$, as would be expected for simple gravitational collapse. This again suggests the absence of important physics from the models. One well-discussed possibility is that of pre-heating (Kaiser 1991; Evrard \\& Henry 1991) in which energy sources such as stars, supernovae and AGN pre-heat the gas in the subclumps that merge to form present-day clusters. More recently, Cavaliere \\etal (1997) have shown that accounting for the effects of shocks in cluster formation (with plausible levels of pre-heating) can lead to $L_{\\rm Bol}\\propto T_{\\rm X}^3$ at intermediate X-ray luminosities, although at the highest luminosities the relation should flatten to $L_{\\rm Bol}\\propto T_{\\rm X}^2$. In this Letter we use the results from multiphase analyses of ASCA spectra and ROSAT High Resolution Imager (HRI) data, for a sample of thirty of the most X-ray luminous known clusters (with $L_{\\rm Bol} > 10^{45}$ \\ergps), to examine the effects of cooling flows on the $L_{\\rm Bol}-T_{\\rm X}$ relation. We show that explicitly accounting for the effects of cooling flows leads to a significant flattening of the slope of the relation, to a value $L_{\\rm Bol} \\approxpropto T_{\\rm X}^2$, in agreement with the theoretical predictions. Throughout this Letter, we assume $H_0$=50 \\kmpspMpc, $\\Omega = 1$ and $\\Lambda = 0$. ", + "conclusions": "" + }, + "9802/astro-ph9802168_arXiv.txt": { + "abstract": "In this paper we investigate whether a misestimate of proper motions could have been a source of substantial systematic errors in the statistical parallax determination of the absolute magnitude of RR Lyrae stars. In an earlier paper, we showed that the statistical parallax method is extremely robust and rather insensitive to various systematic effects. The main potential problem with this method would therefore arise from systematically bad observational inputs, primarily radial velocities and proper motions. In that paper, we demonstrated that the radial velocities have not been systematically misestimated. Here we turn our attention to proper motions. We compare three different catalogs of proper motions --- Lick, Hipparcos and the one compiled by Wan et al.\\ (WMJ). We find that the WMJ catalog is too heterogeneous to be a reliable source. We analyze the sample of 165 halo RR Lyrae stars with either Lick or Hipparcos proper motions. For the stars with both Lick and Hipparcos proper motions we use the weighted means of reported values. Various possible biases are investigated through vigorous Monte Carlo simulations and we evaluate small corrections due to Malmquist bias, anisotropic positions of the stars on the sky, and non-Gaussian distribution of stellar velocities. The mean RR Lyrae absolute magnitude is $M_V=0.74\\pm 0.12$ at the mean metallicity of the sample $\\left<\\rm [Fe/H]\\right>=-1.60$, only $0.01$ mag brighter than the value obtained in the previous study which did not incorporate Hipparcos proper motions. The faint absolute magnitudes of RR Lyrae stars confirmed by this analysis gives strong support to the short distance scale. ", + "introduction": "Popowski \\& Gould (1998, Paper I) showed that the statistical parallax method, which is one of the main luminosity calibration methods for RR Lyrae stars, is extremely robust and is insensitive to several different categories of systematic effects. They proved that the statistical errors are dominated by the size of a stellar sample and that therefore future low-error measurements will have almost no influence on the precision of the estimate of the absolute magnitude, unless the number of sample stars is increased. Consequently, the main immediate avenue to improve the determination of the absolute magnitude of RR Lyrae stars is to eliminate measurement-related systematic errors. There are three main observational pillars on which the statistical parallax method is founded: \\begin{enumerate} \\item radial velocities, \\item proper motions, \\item dereddened apparent magnitudes. \\end{enumerate} Substantial systematics in any of these will result in a miscalibration of the RR Lyrae absolute magnitude. This in turn will affect an RR Lyrae-based distance determination to the Large Magellanic Cloud (LMC) and so will have a crucial impact on the extragalactic distance scale. The method presently known as ``statistical parallax'' is a combination of secular parallax and classical statistical parallax. Secular parallax is based on forcing equality between the three first moments of the velocity distribution (the bulk motion $\\bf w$) as determined from radial velocity and proper motion measurements, while classical statistical parallax is based on forcing equality of the six second moments (the six independent components of the velocity covariance matrix $C_{i j}$). In the modern combined version of statistical parallax one simply determines ten parameters simultaneously by applying maximum likelihood. The ten parameters are an overall distance scaling factor $\\eta$ (relative to an initial arbitrary distance scale) plus the nine first and second moments, $\\bf w$ and $C_{i j}$. In Paper I we analyzed two samples of halo RR Lyrae stars: a kinematically selected sample composed of 162 RR Lyrae stars taken from Layden et al. (1996), and a semi-independent non-kinematically selected sample of stars with metallicities at the [Fe/H]$\\leq -1.5$. The later includes 106 RR Lyrae stars from Layden et al. (1996) and 724 non-RR Lyrae stars from Beers \\& Sommer-Larsen (1995). The Beers \\& Sommer-Larsen (1995) sample was used as a source of additional radial velocities and thus allowed the investigation of possible errors that might be caused by systematic mismeasurement of radial velocities in the Layden et al. (1996) sample. The conclusion was that radial velocities are not a likely source of a systematic error. In this paper we turn our attention to proper motions and establish the reliability of different proper motion catalogs. In \\S 2 we compare the Lick catalog, the Hipparcos catalog, and the catalog compiled by Wan et al.\\ (1980,WMJ) and conclude that the WMJ catalog is too heterogeneous to be a reliable source. We develop uniform procedures to select data and to estimate measurement errors for different subclasses of stars, and we discuss some individual cases. Finally, in \\S 3 we present the result of the statistical parallax method as applied to our new sample and the new sample supplemented by Beers \\& Sommer-Larsen (1995) radial velocities. Our main result is to confirm the previously established absolute magnitude of RR Lyrae stars, which is substantially fainter than the estimates derived using several other methods. ", + "conclusions": "This paper was in part motivated by the recent {\\it Letter} by Tsujimoto et al. (1998) who analyzed a smaller RR Lyrae sample for which there exist Hipparcos measurements. The central value of absolute magnitude obtained by Tsujimoto et al., $M_V = 0.69 \\pm 0.10$ at the average metallicity of $\\left<\\rm [Fe/H]\\right> = -1.58$, is in relatively good agreement with the value derived here, but the two error estimates are in striking disagreement. We would like to stress that the error intrinsic to statistical parallax is completely dominated by sample-size effects and that the high precision of the Hipparcos proper motions allows only a marginal improvement in the precision of the absolute magnitude measurement. The errors quoted by Tsujimoto et al. (1998) for 99 stars are substantially smaller than the theoretical minimum which we derived analytically in Paper I. (We also attempted to confirm Tsujimoto et al.'s (1998) finding of a significant rotation between the Lick and Hipparcos frames. However, we found instead that the rotation is consistent with zero.) The apparent magnitudes of RR Lyrae stars remain the only major unchecked-for-systematics ingredient of the statistical parallax analysis. In appendix A, we describe the results of our (not very conclusive) attempt to address this issue by comparing Hipparcos and Layden (1994) apparent magnitudes. Our analysis suggests a possible correction toward fainter values of the absolute magnitudes of the field RR Lyrae stars, not in the right direction to reconcile the discrepant distance scales. In Paper I we showed that systematic problems with radial velocities are not likely to cause a miscalibration of the RR Lyrae absolute magnitude. Here we have shown that systematic problems with proper motions are also not likely to cause a miscalibration. Because the statistical parallax method is very insensitive to other systematic effects, we conclude that the statistical error of 0.12 mag reported above for the absolute magnitudes of field RR Lyrae stars is the true error. The statistical parallax calibration $M_V=0.74\\pm 0.12$ at $\\left<\\rm [Fe/H]\\right> = -1.60$ is therefore inconsistent at the 2$\\sigma$ level with that derived from main-sequence fitting by Reid (1997) ($ M_V \\sim 0.44 \\pm 0.07$ at [Fe/H]=$-1.6$) and Gratton et al. (1997) ($ M_V = 0.47 \\pm 0.04$ at [Fe/H]=$-1.6$) using Hipparcos parallaxes of nearby subdwarfs. That is, Hipparcos proper motions of RR Lyrae stars confirm a short local distance scale. Gratton (1998) suggested that there may be a difference in the luminosity between globular cluster and field RR Lyrae stars. If RR Lyrae stars in globular cluster belonged to different population the problem would certainly be solved, but Catelan (1998) showed that RR Lyrae stars in these two different environments have essentially the same distribution in the period-equilibrium temperature plane, both at the metal-poor and at the metal-rich ends, suggesting very similar luminosities. He also argued that it is difficult to adjust helium abundance of RR Lyrae stars in different environments in such a way as to produce the required luminosity difference of 0.2 mag and still maintain the striking overlap in the period-equilibrium temperature plane. It is however conceivable that some characteristics of RR Lyrae stars may be sensitive to the presence (globular clusters) or absence (field) of non-canonical deep mixing that occurs during the red giant phase (Kraft et al. 1997)." + }, + "9802/astro-ph9802112_arXiv.txt": { + "abstract": "The phase-transition induced collapse of a neutron star to a more compact configuration (typically a ``strange'' star) and the subsequent core bounce is often invoked as a model for gamma-ray bursts. We present the results of numerical simulations of this kind of event using realistic neutrino physics and a high density equation of state. The nature of the collapse itself is represented by the arbitrary motion of a piston deep within the star, but if any shock is to develop, the transition, or at least its final stages, must occur in less than a sonic time. Fine surface zoning is employed to adequately represent the acceleration of the shock to relativistic speeds and to determine the amount and energy of the ejecta. We find that these explosions are far too baryon-rich ($M_{\\rm ejecta} \\sim 0.01 M_{\\odot}$) and have much too low an energy to explain gamma-ray bursts. The total energy of the ejecta having relativistic $\\Gamma \\gtrsim 40$ is less than $10^{46}$erg even in our most optimistic models (deep bounce, no neutrino losses or photodisintegration). However, the total energy of all the ejecta, mostly mildly relativistic, is $\\sim 10^{51}$erg and, if they occur, these events might be observed. They would also contribute to Galactic nucleosynthesis, especially the $r$-process, even though the most energetic layers are composed of helium and nucleons, not heavy elements. ", + "introduction": "A major goal in modern model building for cosmological gamma-ray bursts is finding a source which provides both the high energies ($\\gtrsim 10^{51}$erg for symmetric explosions) and high Lorentz factors ($\\Gamma \\gtrsim 100$) required to explain the observations (e.g., M\\'esz\\'aros \\& Rees 1993). One often proposed source for this energy is an explosion resulting from the phase transition of a neutron star to a ``strange'' or ``hybrid'' star (Ramaty et al. 1980; Ramaty, Lingenfelter \\& Bussard 1981; Brecher 1982; Ellison \\& Kazanas 1983; Bonazzola 1986; Michel 1988; Haensel, Paczynski \\& Amsterdaamski 1991; Ma \\& Xie 1996; Ma \\& Luo 1996; Shaviv \\& Dar 1996, Qin et al. 1997). Though there is some variation in the models - sometimes a critical mass is achieved for an accreting neutron star, sometimes only a single star is involved - all take note of the large gravitational binding energy of a neutron star and speculate that some fraction of this can be tapped and converted into an outgoing shock wave when the inner core abruptly makes the transition to a more compact state. Because the density is so high, neutrino losses are small, except very near the surface; similarly photodisintegration losses are negligible, so the ``prompt shock'' mechanism that fails to give mass ejection in standard supernova models (e.g. Bethe 1990), might, in this case, deliver large amounts of momentum to the surface layers. Models for strange or ``hybrid'' stars predict central core densities nearly 5 times greater and radii 10-20\\% smaller than neutron stars of the same mass (Rosenhauer, Staubo, \\& Csernai 1991). A rough estimate of the potential energy released, $E \\sim \\frac{G M^2}{R} \\frac{\\Delta R}{R}$, predicts energies on the order of $10^{52}$erg, easily sufficient for gamma-ray burst models (e.g., Ma \\& Xie 1996). However, to get the required high Lorentz factors ($\\Gamma$), this energy must be concentrated into a thin layer near the neutron star surface so that only a small amount of mass is ejected.\\footnote{To reach a mean $\\Gamma$ of 100, a $2 \\times 10^{51}$erg explosion must eject less than $10^{-5} M_{\\odot}$.} Competing with this sink for the gravitational energy is an increase in the internal energy that occurs when matter moves to higher gravitational potential, so the shock does not carry an energy equal to the entire change in gravitational potential. And of course the matter ejected will have a distribution of kinetic energies with most of the mass concentrated at low energies. In this paper, we show that while the phase-transitions of neutron stars can indeed impart total energies to their ejecta in excess of $10^{50}$erg, they fail to deposit enough energy at high $\\Gamma$s to explain gamma-ray bursts. The phase-transition induced collapse of neutron stars is therefore {\\it not} a viable gamma-ray burst mechanism. ", + "conclusions": "" + }, + "9802/astro-ph9802324_arXiv.txt": { + "abstract": "GRO J2058+42, a transient 198 second x-ray pulsar, was discovered by the Burst and Transient Source Experiment (BATSE) on the {\\em Compton Gamma-Ray Observatory (CGRO)}, during a ``giant\" outburst in 1995 September-October. The total flux peaked at about 300 mCrab (20-50 keV) as measured by Earth occultation. The pulse period decreased from 198 s to 196 s during the 46-day outburst. The pulse shape evolved over the course of the outburst and exhibited energy dependent variations. BATSE observed five additional weak outbursts from GRO J2058+42, each with two week duration and peak pulsed flux of about 15 mCrab (20-50 keV), that were spaced by about 110 days. An observation of the 1996 November outburst by the {\\em Rossi X-ray Timing Explorer (RXTE)} Proportional Counter Array (PCA) localized the source to within a 4\\arcmin\\ radius error circle (90\\% confidence) centered on R.A. = 20$^h$ 59$^m$.0, Decl. = 41\\arcdeg 43\\arcmin\\ (J2000). Additional shorter outbursts with peak pulsed fluxes of about 8 mCrab were detected by BATSE halfway between the first four 15 mCrab outbursts. The {\\em RXTE} All-Sky Monitor detected all 8 weak outbursts with approximately equal durations and intensities. GRO J2058+42 is most likely a Be/X-ray binary that appears to outburst at periastron and apastron. No optical counterpart has been identified to date and no x-ray source was present in the error circle in archival {\\em ROSAT} observations. ", + "introduction": "In the last 25 years more than 40 accretion-powered x-ray pulsars have been detected. About half of these are transient, of which 12 have known Be star companions. Neutron stars with Be companions accrete material from the slow, dense, stellar outflow thought to be confined to the equatorial plane of the Be star. Recent long term studies by the Burst and Transient Source Experiment (BATSE) on the {\\em Compton Gamma-Ray Observatory (CGRO)} revealed that Be/X-ray binaries exhibit series of often periodic outbursts. These outbursts are sometimes associated with ``giant'' outbursts accompanied by high spin-up rates and luminosities. BATSE observed four additional accreting x-ray pulsars, without identified companions, which are believed to also be Be/X-ray binaries because their temporal behavior closely resembles that of systems with Be companions (\\cite{Bildsten97}). This paper reports the discovery and temporal behavior of the fourth member of this group, GRO J2058+42. A 198 second periodic signal was observed in the BATSE data starting on 1995 September 14. At the same time a new source was also detected by Earth occultation measurements, which measure phase-averaged (total) flux. A location was determined from both the pulsed data and the non-pulsed data with a 95\\% confidence error box of about 4\\arcdeg\\ $\\times$ 1\\arcdeg\\ (\\cite{Wilson95}). A {\\em CGRO} target of opportunity was declared and the spacecraft was reoriented to allow scans of the region by the Oriented Scintillation Spectroscopy Experiment (OSSE), resulting in an improved 30\\arcmin\\ $\\times$ 60\\arcmin\\ (95\\% confidence) position (\\cite{Grove95}). The total flux peaked at about 300 mCrab (20-100 keV) and the pulsed flux (RMS deviation from mean) peaked at 140 mCrab (20-50 keV) on 1995 September 27. This bright outburst continued until 1995 October 30. A search of archival BATSE data from 1991 April until this bright outburst showed no previous outbursts. An analysis of BATSE data following the bright outburst initially revealed three much weaker outbursts each lasting about two weeks with pulsed flux peaking at 15-20 mCrab (20-50 keV). These outbursts were spaced by about 110 days which allowed the peak of the next outburst to be predicted. A target of opportunity scan of the OSSE/BATSE error box was performed on 1996 November 28 with the {\\em RXTE} Proportional Counter Array (PCA) yielding a 90\\% confidence 4\\arcmin\\ radius error circle centered on R.A. = 20$^h$ 59$^m$.0, Decl. = +41\\arcdeg 43\\arcmin\\ (\\cite{Wilson96}). BATSE also detected the source from 1996 November 23- December 1. Another outburst was detected by BATSE about 110 days later (1997 March 16-20.) Three shorter outbursts with peak pulsed fluxes of about 8 mCrab were detected by BATSE halfway between the first four 15-20 mCrab outbursts. The {\\em RXTE} All-Sky Monitor detected all 8 weak outbursts with approximately equal durations and intensities. An archival search of {\\em ROSAT} data found no sources within the error circle (J. Greiner, 1997, private communication.) No optical counterpart has been found to date. In this paper we present the BATSE and {\\em RXTE} observations of GRO J2058+42. Our observations with BATSE include histories of pulse frequency and pulsed flux from 1995 September - 1997 March, a history of phase-averaged flux for the ``giant\" outburst (1995 September-October), and pulse profile variations dependent upon energy and outburst phase of the ``giant\" outburst. An {\\em RXTE} PCA observation of a weak outburst on 1996 November 28 includes a fit to the scan data used to better locate GRO J2058+42 and pulse profiles. We compare BATSE and {\\em RXTE} ASM observations of the 8 weak outbursts from 1995 December - 1997 March. We then discuss the implications of our results. ", + "conclusions": "The optically identified companions for transient pulsars with pulse periods longer than one second are all Be or Oe stars. Long term studies of these Be/X-ray binaries have demonstrated that giant outbursts followed by or interspersed with a series of periodic normal outbursts appear to be typical behavior for Be/X-ray binaries (\\cite{Bildsten97}). GRO J2058+42 is most likely a Be/X-ray pulsar because it exhibits transient outbursts recurring with a presumed orbital period of 110 days and it has both giant and normal outbursts. An orbital period of 110 days places GRO J2058+42 along the orbital period spin period correlation for Be/X-ray binaries (\\cite{Corbet86}, \\cite{Waters89a}.) The giant outburst of GRO J2058+42 showed enough dynamic range that we can use the relationship between torque and observed flux to test accretion theory. Simple accretion theory assumes that material from the companion star is flowing onto a rotating neutron star with a magnetic field. The magnetic field is so strong that it determines the motion of the material in a region of space surrounding the neutron star called the magnetosphere. The size of this region, the magnetospheric radius $r_{\\rm m}$, is defined to be the distance from the neutron star at which all magnetic field lines are just closed loops. Another important length scale is the corotation radius, $r_{\\rm co}$, the distance from the neutron star where centrifugal forces just balance local gravity. If $r_{\\rm m} > r_{\\rm co}$, accretion is not expected to occur. Accretion theory predicts $r_{\\rm m} \\propto \\dot M^{-2/7}$ for $r_{\\rm m} < r_{\\rm co}$ for for disk or wind accretion. The accretion torque $N$ is given by $N \\propto \\dot M \\sqrt r_{\\rm m}$. We can assume the bolometric flux $F_{\\rm bol}$ is related to the mass accretion rate $\\dot M$ by $\\dot M \\propto F_{\\rm bol}$. Therefore, simple accretion theory predicts $\\dot \\nu \\propto F_{\\rm bol}^{6/7}$, where $\\dot \\nu$ is the spin-up rate of the neutron star (\\cite{Henrichs83} and references therein). Figure~\\ref{fig:flvsfd} shows the observed 20-100 keV flux as measured by Earth occultation plotted versus the measured frequency derivative $\\dot \\nu$ during the giant outburst. The frequency derivatives were generated by an search performed on the BATSE DISCLA 20-50 keV data over a grid of trial frequencies and frequency derivatives. This technique was identical to that described in section~\\ref{sec:giant} except that the pulse profiles were shifted according to a grid of frequency offsets and frequency derivatives, rather than only frequency offsets. The search yielded maximum spin-up rates of 5 $\\times 10^{-12}$ Hz s$^{-1}$ in the weak outbursts. If we assume that the orbital contribution was small compared to the intrinsic $\\dot \\nu$, we can treat the measured $\\dot \\nu$ as equal to the intrinsic $\\dot \\nu$ and compare it to the total flux in the BATSE energy range. Clearly, $\\dot \\nu$ and the total flux were correlated, hence an accretion disk was likely to be present. The power law $\\dot \\nu \\propto F^{\\gamma}_{\\rm obs}$ with $\\gamma = 6/7$, which is predicted by accretion torque theory, is shown in figure~\\ref{fig:flvsfd}. Also shown is the best fit curve with $\\gamma \\simeq 1.2$. Orbital contributions to the torque and the fact that BATSE does not measure the bolometric flux may explain the discrepancy between the measured slope and accretion theory. Interestingly, a similar fit to EXOSAT data from the 1985 giant outburst of the Be/X-ray binary EXO 2030+375 also gave $\\gamma \\simeq 1.2$ (\\cite{Reynolds96}). In addition to testing accretion theory, we can use the maximum $\\dot \\nu$ ($\\dot \\nu_{\\rm max} = 2.48 \\times 10^{-11}$ Hz s$^{-1}$) from the giant outburst to estimate a distance to GRO J2058+42. The luminosity is related to $\\dot \\nu$ by \\begin{equation} L_{37} = 4.83 \\times 10^{13} \\mu_{30}^{-1/3} m^{1/2} R_6^{-1} I_{45}^{7/6} n(\\omega_s)^{-7/6} \\dot \\nu^{7/6} \\end{equation} where the notation is $\\mu_{30} = \\mu/10^{30}$ G cm$^3$, $ m = M_x/M_{\\sun}$, $R_6 = R/10^6$ cm, $I_{45} = I/10^{45}$ g cm$^2$, and $L_{37} = L_{\\rm acc}/10^{37}$ ergs s$^{-1}$. The neutron star has magnetic moment $\\mu$, radius R, mass M$_x$, moment of inertia I, accretion luminosity L$_{\\rm acc}$, and ``dimensionless torque\" $n(\\omega_s)$. We assumed the very slow rotator case of $n(\\omega_s) \\approx 1.4$ (\\cite{Henrichs83} and references therein). We used typical values (m = 1.4, $R_6 = 1$, $I_{45} = 1$). Since $\\mu_{30}$ is less well known, we calculated bolometric luminosities, L$_{\\rm acc} \\simeq 0.7-3.5 \\times 10^{38}$ ergs s$^{-1}$, for the range $\\mu_{30} = 0.1 - 10$. These luminosities were comparable to the Eddington limit for neutron stars, L$_{\\rm Edd} \\simeq 2 \\times 10^{38}$ ergs s$^{-1}$. Next, to calculate the distance, we can use the 20-100 keV flux corresponding to $\\dot \\nu_{\\rm max}$, $F_{\\rm max} = 4.76\\times10^{-9}$ ergs cm$^{-2}$ s$^{-1}$. Lastly, we must calculate a bolometric correction $\\alpha$, the ratio between the BATSE flux and the total 2-100 keV flux. Simultaneous BATSE and {\\em RXTE} measurements from 1996 November yielded a bolometric correction of $\\alpha \\simeq 0.4$. Then the distance $d$ is given by \\begin{equation} d = \\frac{1}{3.09\\times10^{21} \\mathrm{cm}} (\\frac{\\alpha L_{\\rm acc}}{4\\pi F_{\\rm max}})^{1/2} \\mathrm{kpc} \\end{equation} which gives distances of 7-16 kpc for GRO J2058+42. For these distances the estimated bolometric luminosity for the 1996 November outburst is L$_{\\rm bol} \\simeq 0.2-1 \\times 10^{37}$ ergs s$^{-1}$. Bolometric luminosities for the intermediate weaker outbursts (in the BATSE data) appear to be similar to the 1996 November outburst, but are more uncertain because complete spectral and pulsed fraction information are unavailable. The weak outbursts provide interesting, but difficult to interpret, information about this system. The BATSE data indicate brighter and longer outbursts every $\\approx$110 days with weaker outbursts in between, while the ASM data show outbursts every 54 days that are similar in intensity and duration (fig~\\ref{fig:xte} \\& \\cite{Corbet97}). The intermediate weaker outbursts in the BATSE data must be of a different character than the stronger outbursts. The fact that BATSE detects a 110 day periodicity and {\\em RXTE} detects a periodicity at half that period implies that either a spectral change or a change in the pulse fraction is occurring every other outburst, indicating a difference in accretion mode. An analysis of the three {\\em RXTE} ASM energy bands shows that both sets of outbursts have comparable intensity in the 5 - 12 keV band. This implies that if a spectral change is occurring it is above 12 keV. The observed periodicity may be interpreted in two ways: (1) the 110 day cycle observed in the BATSE data is the orbital period and 2 outbursts are occurring each orbit; or (2) the orbital period is 54 days and the spectrum (or pulse fraction) alternates every other orbit for at least 9 orbits. We can find no plausible explanation for an alternating change in spectrum or pulse fraction that persists for at least 9 orbits if the outbursts are produced by periastron passage in a 54 day orbit. Therefore we interpret the observed 110 day periodicity as the orbital period. We propose that GRO J2058+42 is undergoing periastron and apastron outbursts in a 110 day orbit. A neutron star in an inclined orbit combined with concentration of material in the equatorial plane of the Be star is a possible explanation for the outburst behavior observed in GRO J2058+42. The same velocity kicks that produce an eccentric orbit could also produce an inclined orbit. Outbursts would occur twice per orbit when the neutron star crosses the equatorial plane of the companion (\\cite{Priedhorsky87}). From the BATSE data, we measured the separation between the brighter outbursts to be 110 $\\pm$ 3 days and the separation between brighter and weaker outbursts to be $\\Delta T = 54.5 \\pm$ 2.3 days. If we assume the separation between the brighter outbursts is the orbital period, $P_{\\rm orb}$, then $\\Delta T/P_{\\rm orb}$ falls in the 95\\% confidence interval $0.45 < \\Delta T/P_{\\rm orb} < 0.55$. For this interval the separation between periastron and the line where the neutron star's orbit intersects the equatorial plane of the companion is $\\nu <$ 12.1\\arcdeg\\ for a typical eccentricity, $e = 0.35$. Hence this mechanism would produce outbursts near periastron and near apastron. This mechanism does not, however, explain the intensity differences seen only in the higher energies (20-50 keV). Although an inclined orbit has often been proposed as a mechanism for outbursts in Be/X-ray binaries, two outbursts per orbit have not been seen in other Be/X-ray binaries to date. However, such behavior has been observed in wind-fed systems such as GX 301-2 (\\cite{Koh97}). Another explanation which does not require an inclined orbit suggests that two different accretion mechanisms could be at work in this system. An accretion disk is likely to be present during giant outbursts. Its presence helps to explain the large and steady spin-up rates seen during giant outbursts. A peak spin-up rate of 2.5 $\\times 10^{-11}$ Hz s$^{-1}$ was observed in GRO J2058+42. Bildsten et al. (1997) suggest that an accretion disk could remain after a giant outburst. Then normal outbursts could be produced by large tidal torques on the accretion disk during periastron passage. This could explain the series of normal outbursts following the giant outburst. The apastron outbursts could be produced by accretion from the equatorial outflow from the Be star. A slow dense outflow could produce outbursts at apastron (\\cite{Waters89b}). The two outbursts per orbit would be produced by different accretion mechanisms which could produce differences in spectra or in pulse fraction. Clearly, additional observations are needed to fully understand this source." + }, + "9802/astro-ph9802144_arXiv.txt": { + "abstract": "We present near-infrared (NIR) surface photometry of a sample of 14 early-type spirals with observed rotation curves. In this first paper, we report the results of two-dimensional parametric and non-parametric decompositions to separate the bulge and disk components; the parametric bulge is modeled with a generalized exponential law of integer index $n$, and the disk with a simple exponential. We find that the derived bulge parameters, for a given galaxy, vary systematically with the bulge shape index $n$. The mean early-type bulge has a best-fit $n$~=~2.6, and 80\\% of the sample has best $n$ of 2 or 3. Bulges are rarely spherical; the median bulge intrinsic ellipticity is 0.33. The median early-type disk has $(J-K)_d$ more than 0.1~mag bluer than the bulge, and a NIR disk surface brightness more than 1 mag~arcsec$^{-2}$ brighter than later-type disks. Our data are consistent with the well-established correlation of both bulge and disk surface brightness with physical scale length, and we note that the location of bulges within this projection of the fundamental plane depends on their shape index $n$. In agreement with previous work, the ratios of bulge and disk scale lengths are consistent with a constant value $r_e/r_d$~=~0.3; however, such value again depends on the bulge index $n$, implying that claims for a scale-free Hubble sequence may be premature. ", + "introduction": "Surface brightness distributions of external galaxies have been studied for many years but a reliable decomposition into structural components is often difficult to obtain. The reliability of models and techniques has been questioned by several authors (Kent \\cite{kent:1986}; Schombert \\& Bothun \\cite{schombert}; Byun \\& Freeman \\cite{byun:freeman}), and there are several points that need further study and refinement. First, brightness distributions have often been studied as one-dimensional (1-D) radial profiles extracted by averaging along elliptical annuli; such ellipses can deviate considerably from the actual isophotes, especially in highly inclined systems with a luminous bulge. The result provides therefore a distorted profile of the surface brightness along the major axis. Second, the seeing must be properly taken into account (c.f., Schombert \\& Bothun \\cite{schombert}), especially when studying the central regions where brightness gradients are highest. Third, the choice among the parametric forms of brightness distribution (exponentials, Hubble and de Vaucouleurs laws, and so on) does not depend on any physical argument, but only on their ability to fit the data. Finally, the effects of internal extinction should be taken into account. The new near-infrared (NIR) panoramic detectors have made imaging in the 1 to 2.5\\micron\\ regime rather straightforward with sensitivity and accuracy comparable to those attainable at optical wavelengths. The NIR bandpasses have been advocated to be the ideal ones to study the characteristics of the galactic backbones, that is the stellar populations which make up the mass distribution in a galaxy (e.g., Rix \\& Rieke \\cite{rix:rieke}). The reason for this is twofold: first, the extinction is lower, by a factor of ten between the $B$ and $K$ bandpasses; and second, the emission of old stellar populations peaks in the NIR. In this context, we have undertaken a program of NIR imaging of bright spiral galaxies with measured rotation curves. These images are used to decompose the luminosity distribution into bulge and disk components and then to analyze their contribution to the observed rotation curves. To overcome the methodological drawbacks mentioned above, we have developed a parametric technique to fit a two-dimensional (2-D) bulge $+$ disk distribution to the entire image, which takes into account the effect of seeing. In addition, generalizing Kent's non-parametric approach, we have developed an iterative algorithm to deduce from the 2-D brightness distribution the contributions of bulge and disk, again taking into account the seeing. In this first paper, we report the results of the structural decomposition for a restricted sample of early-type spirals; a subsequent paper will describe the inferred mass distributions. Such systems appear well suited for studies of this kind, mainly because of the relative smoothness of their disks and their (reputedly) lower internal extinction. On the other hand, due to the low gas content, rotation in Sa's is not measured as far out as for later types, and hence they are not quite as effective for exploring dark matter properties. In several cases it has been noted that the presence of a dominant bulge might severely complicate the picture and the analysis (Kent \\cite{kent:1988}); these are galaxies with luminous bulges and slowly rising rotation curves, and some of them are also included in our sample. ", + "conclusions": "} We have decomposed $J-$ and $K-$band images of 14 early-type spirals into bulge and disk components. 2D non-parametric solutions and results from fitting a parametric model of generalized exponential $1/n$ bulges and simple exponential disks are compared, and general characteristics of early-type spiral bulges and disks are examined. We find that: \\begin{enumerate} \\item Even using objective and refined techniques, the decomposition in structural components is far from being a robust and unique process. For the parametric methods, significantly different decompositions are obtained for different bulge distribution laws. Non-parametric techniques, on the other hand, appear to be affected by the choice of the ellipticities of the components which are difficult to evaluate objectively. Of the two components, the bulge is the most subject to errors, since the inner part is masked by seeing, and the outer regions are buried beneath the disk. \\item Bulge structural parameters are strongly influenced by the form of the function used to derive them. The same bulge, when fitted with small $n$, appears to be ``denser'' (brighter $\\mu_e$), more compact (smaller $r_e$), and less luminous than when fitted with large $n$. The dispersion of the fitted parameters also increases with $n$. \\item The median early-type bulge has a shape index $n$ between 2 and 3, $\\mu_e^c(K)$=16.8\\magarc , and $r_e$=1.6 kpc. It is also red, with $(J-K)_b$~=~1.06, and redder bulges tend to be ``denser'', that is with brighter $\\mu_e$. \\item As noted by Kent (\\cite {kent:1988}), bulges are rarely spherical. The median intrinsic ellipticity is 0.34, equivalent to a disk with 50$^\\circ$ declination. This restricts the applicability of non-parametric techniques to rather inclined systems and suggests treating with caution the studies which assume spherical bulges. \\item The median early-type disk with $\\mu_d^c(K)$~=~17.1 is more than 1~\\magarc\\ brighter than later-type disks, and bluer than the bulge in $(J-K)$ by more than 0.1~mag. Disk scale lengths agree fairly well with those found by other at different wavelengths, and we confirm a tendency for NIR disk scale lengths to be smaller than those at optical wavelengths (e.g., Peletier et al. \\cite{peletier}). $r_d/R_{25}$ is approximately constant, 0.24, similar to the value of 0.25 found for late-type spiral disks (Giovanardi \\& Hunt \\cite{giova:1988}; Giovanelli et al. \\cite{giovanelli}). \\item Both bulge and disk surface brightnesses correlate with respective scale lengths, consistently with the projection of the fundamental plane for ellipticals and spiral bulges (e.g., Andredakis et al. \\cite{andredakis:peletier}). We note that uncertainties in the decomposition, especially in the shape index $n$, strongly influence the position of a bulge within the FP. Disks appear to reside in a region of this FP projection which is roughly contiguous to that of bulges, extending the correlation to larger radii and fainter surface brightnesses. \\item We confirm the tendency for the ratio of bulge and disk scale lengths $r_e/r_d$ to be constant, noted by de Jong (\\cite{dejong:3}) and Courteau et al. (\\cite{courteau}). However, we find a mean (best-$n$) value $r_e/r_d$~=~0.3, significantly larger than the value found by de Jong and Courteau et al.; our $n$~=~1 value of $r_e/r_d$ of 0.2 agrees roughly with their value of 0.13--0.14, while our $n$~=~4 value is 0.7, more than a factor of 3 larger. We attribute such differences to different bulge parameterizations and caution that if best $n$ varies with morphological type, as suggested by Andredakis et al. (\\cite{andredakis:peletier}), $r_e/r_d$ may not be constant with morphological type, and thus the Hubble sequence may not be scale free as proposed by Courteau et al. (\\cite{courteau}). \\end{enumerate}" + }, + "9802/astro-ph9802002_arXiv.txt": { + "abstract": "Because of their simplicity, axisymmetric mass distributions are often used to model gravitational lenses. Since galaxies are usually observed to have elliptical light distributions, mass distributions with elliptical density contours offer more general and realistic lens models. They are difficult to use, however, since previous studies have shown that the deflection angle (and magnification) in this case can only be obtained by rather expensive numerical integrations. We present a family of lens models for which the deflection can be calculated to high relative accuracy $(10^{-5})$ with a greatly reduced numerical effort, for small and large ellipticity alike. This makes it easier to use these distributions for modelling individual lenses as well as for applications requiring larger computing times, such as statistical lensing studies. A program implementing this method can be obtained from the author\\footnote{or at \\verb+http://www.sns.ias.edu/\\~{}barkana/ellip.html+}. ", + "introduction": "The types of density profiles used to model gravitational lenses have been motivated by observations of lenses in addition to practical considerations. Observed features of galaxies and clusters that can be incorporated into lens models include ellipticity and radially decreasing mass density profiles. But it is essential to be able numerically to calculate quickly the deflection angle and magnification of a light ray due to a lens model, in order to probe the entire range of parameter space when searching for a best fit to an observed lens. Numerical efficiency is even more important in cases with multiple sources such as radio lenses observed at high resolution. Strongly lensed arcs in clusters sometimes lie near cluster galaxies and thus modelling all the lensed features in a cluster may require a combination of many individual lenses. Another application which depends on numerical speed is the statistical study of properties of the images of a lens model for comparison with lens surveys. In attempting to construct models that are as realistic as possible, new features must be introduced carefully, since in a single case of multiple lensing there are a small number of constraints requiring a small number of model parameters. Ellipticity adds just two parameters (magnitude and orientation), and it appears to be essential. Galaxies and clusters often appear elliptical, with significant numbers having axis ratios $b/a$ smaller than $0.5$. Axisymmetric lens models are also excluded by many observed lens systems, in which the images are not colinear with the lens position on the sky. Of course, the asymmetry in each case may also be due in part to external shear from other nearby galaxies or from large-scale structure along the line of sight (Bar-Kana 1996, Keeton et al. 1997). However, the use of elliptical mass distributions has not become common practice because the evaluation of the deflection angle and magnification matrix requires some effort. Bourassa et al.\\ (1973) and Bourassa and Kantowski (1975) (with minor corrections by Bray (1984)) introduced a complex formulation of lensing which allows for an elegant expression of the deflection angle due to a homoeoidal elliptical mass distribution. I.e., this is a projected, two dimensional mass distribution whose isodensity contours are concentric ellipses of constant ellipticity and orientation. It can be obtained, e.g., by projecting a three-dimensional homoeoidal mass distribution. The complex integral which gives the deflection angle is in practice difficult to separate into real and imaginary parts. Schramm (1990) used an alternative derivation to obtain the deflection without the use of complex numbers, but still requiring a numerical integral for each component of the deflection angle. Elliptical densities have been used for numerical lens modelling, with the ellipticity allowed to vary in order to fit the data, by Keeton \\& Kochanek (1997). In order to avoid numerical integration, several alternatives have been suggested to exact elliptical mass distributions. Models where the potential is chosen to have elliptical contours rather than the density are easy to use, since the deflection can be obtained immediately as the gradient of the potential. The imaging properties of elliptical potentials have been investigated extensively (Kovner 1987, Blandford \\& Kochanek 1987 and Kochanek \\& Blandford 1987). They become identical to elliptical densities for very small ellpticities and produce similar image configurations even for moderate ellipticty (Kassiola \\& Kovner 1993, hereafter KK93). However, elliptical potentials cannot represent mass distributions with axis ratios $b/a$ smaller than about $0.5$ because the corresponding density contours acquire the artifical feature of a dumbbell shape, and the density can also become negative in some cases (Kochanek \\& Blandford 1987, KK93). To avoid this problem, Schneider \\& Weiss (1991) proposed a numerical method based on a multipole expansion of the mass distribution, but this expansion converges slowly when used with large ellipticity. In this paper we consider a family of projected density profiles that has been used with the approximate approaches to ellipticity discussed above. This is the family of softened power-law profiles, which have a constant density within a core radius and approach a power-law fall-off at large radii. In \\S 2 we introduce our notation for softened power-law elliptical mass distributions (SPEMDs) and for softened power-law elliptical potentials (SPEPs), and illustrate further the limitations of SPEPs. In \\S 3 we present and simplify the quadrature solution of Schramm (1990) for the deflection angle. We show that for the SPEMDs it is possible to approximate the integrand so that the integral can be done analytically. Although the result is a sum of series expansions, each series converges rapidly to high accuracy, even for mass densities with arbitrarily high ellipticity. We show how to similarly evaluate the magnification matrix. We also derive an expression for the gravitational potential, but it cannot be evaluated without a numerical integration. Finally, in \\S4 we summarize our results. ", + "conclusions": "A mass density profile with elliptical isodensity contours is a natural lens model to try when axisymmetric models fail. Previously, however, the easiest way of evaluating the deflection has been to numerically integrate the solutions of Schramm (1990). After simplifying these solutions we have shown that for the family of SPEMD mass distributions, the deflection angle and magnification matrix can be evaluated very fast and with high accuracy. Our implementation achieves a relative accuracy of $5\\times 10^{-6}$ in the deflection and $6\\times 10^{-4}$ in the magnification while running 20 times faster than a procedure based on the numerical integrations. We have also derived an expression for the potential, although this quantity must be numerically integrated. As noted by Schneider \\& Weiss (1991), combinations of two or more SPEMDs with different parameters can be used to construct more general density profiles with several scales. We thus expect SPEMDs to be more widely used, particularly for cases of high ellipticity in which the alternative SPEPs develop the artificial feature of dumbbell-shaped contours." + }, + "9802/astro-ph9802234_arXiv.txt": { + "abstract": "\\noindent From considering the effect of $\\gamma$-$\\gamma$ interactions on recently observed TeV gamma-ray spectra, improved limits are set to the density of extragalactic infrared (IR) photons which are robust and essentially model-independent. The resulting limits are up to two orders of magnitude more restrictive than direct observations in the 0.025-0.3 eV regime. These limits are used to improve constraints on radiative neutrino decay in the mass range above 0.05eV and on Very Massive Objects (VMOs) as providing the dark matter needed to explain galaxy rotation curves. ", + "introduction": "The extragalactic background IR field potentially contains a wealth of information relevant to both cosmology and particle physics (see for example \\cite{Bond}), but has so far eluded any conclusive detection. Direct measurement of the characteristics of this radiation are frustrated by the dominance of local, galactic IR sources. TeV gamma-ray astronomy provides a means to study the IR background indirectly, free of such complications, by looking for modifications to high energy gamma ray spectra due to interactions with this background field via the process $\\gamma \\gamma \\longrightarrow e^{+} e^{-}$. This idea was first noted by Gould and Schr\\'{e}der \\cite{Gould} and has been more recently restated by Stecker and de Jager \\cite{Stecker}. Previous IR studies making use of this phenomenon have generally used spectral data from the active galactic nucleus (AGN) Mrk 421 taken by the EGRET \\cite{Hartman} and Whipple \\cite{Mohanty} experiments at GeV and TeV energies, respectively. These studies have relied on extrapolating flux measurements between these two energy regimes in an effort to determine the inherent shape of the source spectrum \\cite{Stecker} \\cite{Dwek} \\cite{DeJager} \\cite{Biller}. This approach has several difficulties: {\\bf 1)} Since the flux level is known to be highly variable at TeV energies, the relative normalization between EGRET and Whipple flux measurements is not known; {\\bf 2)} The flux of IR photons cannot be assumed to be zero in the intervening energy range; and {\\bf 3)} The shape of the inherent source spectrum is not known and cannot be assumed to necessarily follow a single power law from GeV to TeV energies. Indeed, current best fit values for Mrk 421 at TeV energies indicate a steeper spectrum than is reflected by the EGRET data \\cite{Zweerink}, suggesting that the spectral shape changes in the intervening energy range. In addition, most of the analyses have assumed a particular model to describe the shape of the IR photon spectrum. Even a more recent attempt at a less model-dependent limit using data from Mrk 501 nevertheless assumes a power-law source spectrum, an arbitrary IR flux normalization and constant IR energy densities over energy ranges spanning up to an order of magnitude in extent \\cite{Stanev}. Limits derived in this manner may not apply if the background IR spectrum is of a different shape and, in particular, cannot necessarily be used to constrain mechanisms which might produce sharp features in the spectrum on a scale smaller than that of the assumed model. This current paper will attempt to define more robust limits based on recent high energy observations of AGN which do not rely on extrapolation from lower energies or the assumption of a strict power-law source spectrum and are independent of the IR model down to scales corresponding to a factor of 2 change in IR energy. These limits will then be used to constrain significantly the possible contribution of Very Massive Objects (VMOs) to the dark matter problem and also to place improved bounds on the radiative decay of massive neutrinos in the mass range above $\\sim0.05$ eV, which is in the regime of interest to the atmospheric neutrino problem. ", + "conclusions": "Recent data from very high energy gamma-ray observations of active galaxies has been used to set improved upper limits to the infared background radiation which are robust and relatively model-independent. These limits cast doubt on VMO models as providing the explanation for dark matter in galactic halos and tighten constraints on the radiative decay of relic neutrinos in the mass range near 0.1 eV by approximately two orders of magnitude. Substantial improvement in the ability to study the infrared background via $\\gamma$-$\\gamma$ interactions is expected with the next generation of instruments which will allow the behavior of more distant sources to be examined. In fact, the procedure used to derive the limits given in Table 1 may also be used to determine lower limits to the maximum distance out to which TeV gamma-ray telescopes can be expected to probe. This has been done for the gamma-ray energy range of 0.4-10 TeV, assuming that a source remains ``visible'' out to an optical depth of 2 and that the majority of the IR was produced prior to the epoch of the sources under study. IR and UV upper bounds from direct observations in the energy range greater than 2 eV have also been used to calculate this minimum distance for gamma-ray energies of 0.05 and 0.1 TeV. The results are shown in figure 4. The current generation of ground-based, gamma-ray telescopes have a typical lower energy threshold of $\\sim$0.5 TeV and are thus expected to be able to see sources possessing redshifts up to or beyond $z$=0.1. The next generation of instruments is expected to have an energy threshold in the region of 0.05 to 0.1 TeV, and will therefore be able to see out to a redshift of at least $z$=0.5. This represents a substantially larger volume of space than has so far been explored at TeV energies. This, and the known preponderance of AGN at larger redshifts, indicates a bright future for ground-based, gamma-ray astronomy and strongly suggests a drive towards higher duty cycles, larger apertures and greater dynamic ranges for such instruments to explore this regime thoroughly. This work has been supported in part by PPARC, Forbairt, the US Department of Energy, NASA and the Smithsonian Institution." + }, + "9802/astro-ph9802078_arXiv.txt": { + "abstract": "\\noindent We report on highlight results on celestial sources observed in the high energy band ($>20$~keV) with \\B. In particular we review the spectral properties of sources that belong to different classes of objects, \\ie stellar coronae (Algol), supernova remnants (Cas~A), low mass X--ray binaries (Cygnus X--2 and the X--ray burster GS1826--238), black hole candidates (Cygnus X--1) and Active Galactic Nuclei (Mkn~3). We detect, for the first time, the broad-band spectrum of a stellar corona up to 100 keV; for Cas~A we report upper limits to the $^{44}$Ti line intensities that are lower than those available to date; for Cyg X--2 we report the evidence of a high energy component; we report a clear detection of a broad Fe K line feature from Cyg X--1 in soft state and during its transition to hard state; Mkn~3 is one of several Seyfert 2 galaxies detected with \\B\\ at high energies, for which Compton scattering process is important. ", + "introduction": "High energy properties of celestial X--ray sources give important information to understand their radiation mechanisms and the energetic processes occurring in them and/or in their environments. Hard X--ray emission ($>20$~keV) is currently observed from several classes of X--ray sources. Galactic X--ray sources that are known emitters of hard X--rays include black-hole candidates, X--ray pulsars, weak-magnetic-field neutron stars in Low Mass X--ray Binaries (LMXRBs), mainly X--ray bursters (XRBs), Cataclismic Variables (CV), in particular Polars, Crab-like supernova remnants. High energy spectra of black-hole candidates (BHC) have permitted to infer the presence of Comptonization processes of soft photons occurring close to the black-hole (\\eg in a disk corona). XRBs, that are weakly magnetized neutron stars (with surface field intensity $B\\leq 10^{10}$--$10^{11}$~G), turned out to be hard X--ray emitters, once the sensitivity of the high energy instruments was increased at the 10~mCrab level (see \\cite{tavani97} for a recent review). From their spectral properties, similarities with and diversities from BHCs have been inferred, like the presence of an accretion disk that can extend, as in the case of BHC, close to the surface of the compact object, and the presence of an additional component of soft photons that, unlike in BHCs, originates from the neutron star surface and can be a major source of thermal emission and electron cooling through Comptonization. X--ray pulsars are well known emitters of hard X--rays. Observations in the hard X--ray band are relevant in order to get a measurement of the magnetic field intensity at the neutron star surface. Even if current models of the X--ray spectrum of these objects are still unsatisfactory at high energies, the measurement of cyclotron resonance features gives a direct estimate of the intensity of the neutron star magnetic field strength~\\cite{dalfiume97}. Emission from young shell-like supernova remnants mainly extends to low energies ($<20$~keV). Detection of hard X--rays with determination of their spectral properties can provide important information on the emission mechanism (thermal vs.\\ non thermal, like synchrotron radiation). \\begin{figure} \\centerline{\\epsfig{figure=fig1_new.ps,height=8.5cm,width=\\columnwidth}} \\caption[]{Unified model for AGNs. Adapted from \\cite{urry95}.} \\label{fig:fig_1} \\end{figure} From stellar coronae, apart the Sun, hard X--rays have never been observed. As we will see, this gap has been filled with \\B. Among the extragalactic X--ray sources, hard X--ray emission is observed from Active Galactic Nuclei (AGNs), that include Seyfert galaxies of both types (1 and 2), radio quiet QSOs and radio loud QSOs (which include blazars). A great effort is currently under way to interpret the different classes of AGNs in a unified scheme, which is sketched in Fig.~\\ref{fig:fig_1}. The basic energy production mechanism is accretion of matter onto a massive black hole ($\\approx 10^8$~M$_\\odot$) via an accretion disk. A massive toroid of larger radius (in the range from several parsecs to few tens of parsecs~\\cite{krolik94}) is assumed to surround the accretion disk. Depending on the configuration of the disk with respect to the toroid, on their relative sizes and distances, and on the viewing angle, an AGN should show different observational features and thus fall in one of the different classes above mentioned. The above scheme is being tested also for stellar-mass black holes (\\eg Cygnus X--1). Thus the unified scheme can be a general picture to interpret galactic and extragalactic black holes, accreting matter via an accretion disk. Given many similarities in the X--ray emission from stellar mass BHs with low-magnetic-field neutron stars in LMXRBs, the unified scheme now applied to AGNs could be extended to several classes of X--ray sources. Hard X--ray spectral properties of these sources can provide unique information to diagnose the presence of a black hole versus a weak-magnetic-field neutron star, to test the unified model for AGNs and its validity for stellar mass BHCs. Thank to a broad energy band of operation (0.1--300~keV) and a uniform flux sensitivity in this range, \\B\\ \\cite{BeppoSAX} has the unmatched capability of simultaneously sampling the spectrum of X--ray sources over more than three decades of energy. The SAX/PDS instrument \\cite{PDS}, with a sensitivity of about 1~mCrab at 100~keV, allows an accurate determination of the spectrum of many X--ray sources at the highest energies (13--200~keV). For the brightest sources ($>10$~mCrab), the HPGSPC instrument (6--60 keV) \\cite{HPGSPC} provides the spectral coverage necessary to match the information provided by the low energy instruments LECS (0.1--10~keV) \\cite{LECS} and MECS (2--10~keV) \\cite{MECS} telescopes and the PDS. Here we review some relevant results obtained with \\B\\ during its Performance Verification Phase and first year Core Program, with particular focus on the PDS instrument. The spectral deconvolution was performed with the XSPEC software package, by using the instrument response function distributed from the \\B\\ Scientific Data Center. ", + "conclusions": "Observations of celestial sources with \\B\\ show new key results on their high energy ($>20$~keV) spectral properties. Strong high energy emission has been detected from Algol during 1 day long flare. The emission appears to be the tail of the thermal X--ray radiation. LMXRBs are known sources of low energy X--ray emission, but the high energy emission is well known only for a small part of them (see a recent review by \\cite{tavani97}). We detected for the first time a non thermal high energy component from the Z source Cygnus X--2 and derived, for the first time with the same satellite, the broad--band (0.1--200~keV) photon spectrum of an X--ray burster during a transient hard X--ray outburst. We expect to be able, using the three observations available, to study the relative behaviour of the high energy components with respect to the low-energy one as a function of the high energy flux of GS1826--238. We have reported the clear detection of a broad Fe K line feature from Cygnus X--1 during a soft state of the source and we have discussed its behaviour for different spectral states. The Fe fluorescence emission does not appear to be consistent with a constant disk emission region. As far as the AGNs are concerned, the detection of power laws with strong absorptions at lower energies and reflection components, is now observed for several Seyfert 2 galaxies. We have shown the outstanding example of Mkn~3." + }, + "9802/astro-ph9802287_arXiv.txt": { + "abstract": "We compare the mass functions obtained analytically, in the framework of an extended Press \\& Schechter (PS) formalism, in a previous paper to the numerical mass functions obtained in N -- body simulations, using different algorithms to define objects in the density field. After discussing the properties of the algorithms, we show that the mass function obtained using the friend -- of -- friend algorithm reproduces best the scaling behaviors predicted in the extended PS formalism. Following this statistical analysis, we show that it is possible in the framework of our extended PS formalism to reproduce the mass function but also, and for the first time, the initial statistical properties of structures and their collapse time. This allow to present a ``coherent'' picture of structure formation which can account for the initial, final and dynamical properties of structures. ", + "introduction": "This paper is the second part of a study on the influence of non--linear dynamics on the mass function of cosmic structures. In a companion paper (\\cite{AA97}, thereafter P1) we studied analytically the influence of the shear and of the tide on the mass function in a generalized Press \\& Schechter (1974, thereafter PS) formalism for a critical universe. We first showed that the mass function can be directly related to a given dynamical model through a selection function which gives the probability that a Lagrangian fluid element with a given initial density contrast will be in a collapsed structure at the present epoch. We also emphasized two invariance properties of the mass function in such a formalism, the well-known self-similarity in time, and the possibility of factorizing the mass function as the product of a function which depends only on the dynamical model and another function which depends only on the power spectrum of initial density fluctuations (P1). These time and spectral scaling properties are strong predictions of the generalized PS formalism, they will therefore be used to test the validity of such an approach. In the present paper, we extend our analytical work by a comparative analysis with numerical simulations. These ones have been used by many authors to test the PS approach. The first successful attempt was made by Efstathiou et al. (1988) using however a small number of particles. They used a rather simple and easily tractable algorithm to define collapsed objects in the numerical density fields, namely the Friend -- of -- friend algorithm (thereafter FOF). They conclude that the standard PS mass function reproduces well the numerical data, although the number of objects with mass $M \\simeq M_*$ seems to be overestimated ($M_*$ precedes the cut-off at large mass). Gelb \\& Bertschinger (1994) have presented a new algorithm (DENMAX) which brings together particles around maxima of the density field. They showed that the resulting mass function differs from the FOF mass function. Lacey \\& Cole (1994) compared also the mass functions obtained with different algorithms. They used mainly FOF and another algorithm that detects spherically symmetric density maxima in the field (spherical overdensity, thereafter SO). These works emphasize that the resulting mass function depend strongly on the chosen algorithm to define structures in the density field. A similar conclusion was outlined by Eke et al. (1996) in the case of the correlation function of clusters in a CDM cosmogony. In this paper we work then towards two goals. First, we study as in the previous works, the dependence on different definitions of structures, of the mass functions in the numerical simulations. In particular, we examine if the resulting mass functions satisfy the time and spectral scaling properties, inherent to a PS description of structure formation. Secondly, we investigate the connection between the actual fully non -- linear dynamics of structures in N -- body simulations and the dynamical model used to compute the mass function in analytical approaches. As a matter of fact, even if the number of objects in the field can be relatively well described by, for example, the standard PS mass function, as was argued in previous studies (\\cite{EFWD88}, \\cite{LC93}), this however does not mean that the actual dynamics of each object is correctly described by the spherical model. Finally, we build a complete description of the mass function which allows both a correct statistical description of the number of objects generated in the numerical simulations, but also a right description of their initial statistical properties and of their dynamics. In the next section we present our numerical simulations of scale free power spectra and the different algorithms we used to define objects in the density field. The resulting mass functions and the analysis of their scaling properties are discussed in sect. 3. A complete description, which accounts for both the statistical properties and dynamical origin of the mass function is exhibited in the sect. 4. Finally we discuss our results and conclude in sect. 5. ", + "conclusions": "In P1 we have extended the PS formalism in order to include the effect of the shear and the tide. We have shown than within the PS formalism it was possible to obtain very different behaviors for the mass function which are directly related to the underlying dynamical model. Even though the mass functions are quantitatively very different, they all satisfy, within the PS formalism, two scaling properties (i.e. the time self similarity and the power spectrum scaling (formula (\\ref{spec-scaling})). The aim of this paper was to determine if the PS formalism could explain coherently structure formation both from a statistical and a dynamical point of view. In the first part of this paper we have shown, using numerical simulation, that it was possible to define numerical structures whose mass function obeys both the PS scaling laws. We have tried different algorithms which all lead to different mass functions. The only one which reproduces all the scaling laws is friend--of--friend. However, the resulting universal mass function is still strongly dependent on the percolation length. This dependence of the mass function on the chosen algorithm shows that the choice of the dynamics used in the PS formalism cannot be unique and has to be related to the actual definition of the structures in the density field. In order to have a coherent description of structure formation, the chosen dynamical model should not only reproduce the mass function, but should also give a correct description of the regions selected in the initial density field and of their collapse times. We have shown, using a dynamical prescription proposed in P1, that it is possible to built such a coherent picture. Within this picture, the structures result from a collapse along their third principal axis. The collapse epoch is then a function of the density contrast and of the largest shear eigen-value given by formula (\\ref{3param}) which has three free dynamical parameters. It is then possible to find a set of parameters which reproduces simultaneously the mass function, the initial selection function and gives on average the correct collapse epoch. Such a coherent description is totally impossible with the classical PS formalism even when adjusting the density threshold. This inadequacy comes from the fact that the shear plays an important role in the dynamics of forming halos. This dynamical effect explains the $50\\%$ discrepancy of structures found by FOF-0.2 compared to the classical PS formalism." + }, + "9802/astro-ph9802308_arXiv.txt": { + "abstract": "The energy spectrum of Mkn\\,501 in the TeV energy regime, as measured by the HEGRA (High Energy Gamma Ray Astronomy) \\v{C}erenkov telescopes during its low state in 1995/96 and during a fraction of the 1997 outburst in the TeV energy regime, is shown to place stringent upper limits on the still unknown infrared photon density in the energy region between 3$\\cdot 10^{-3}$ and $3\\cdot 10^{-1}$ eV. Assuming two different shapes for the unknown infrared photon spectrum in this energy range we calculate upper limits on the infrared photon density on the basis of the power-law fit obtained for the observed spectrum up to the maximum energy. ", + "introduction": "The \\hfill cosmologically \\hfill important \\hfill extragalactical \\hfill diffuse \\hfill infrared \\hfill background \\newline (DIRB) in the astronomical window from the optical to the infrared (IR) has not yet been directly determined experimentally due to large sytematic errors driven by local effects. As realized already a long time ago, the detection of extragalactic TeV-$\\gamma$ sources would enable one to indirectly measure this photon density due to unavoidable pair production losses of TeV photons in infrared photon fields (Gould \\& Schr\\'eder 1966, Stecker et al. 1992). Interpreting the Whipple data of Mkn\\,421 (redshift $z$ = 0.031), Stecker et al. (1994) claimed to have seen an exponential cut-off resulting from cosmic absorption of an otherwise smooth power law $\\gamma$-ray spectrum. The authors suggested that the IR-density is given by $n(\\epsilon)\\approx (0.8^{+0.6}_{-0.4})\\cdot 10^{-3}\\epsilon^{-2.6}(h/0.75)$ cm$^{-3}$eV$^{-1}$ where $h$ refers to the normalization factor of the Hubble constant ($H_\\circ = h 100$km s$^{-1}$ Mpc$^{-1}$). A similar analysis has been carried out by Dwek \\& Slavin (Dwek \\& Slavin 1994). The inferred large values of the diffuse near-infrared background density extrapolated to the infrared would imply that it is virtually impossible to discover any extragalactic source above a few TeV. However, Biller et al. (Biller et al. 1995) correctly pointed out that unless the source spectrum is known, one measured TeV source can only yield an upper limit, because the cut-off may be due to internal absorption at the source. They obtain a conservative upper limit of $\\epsilon^2 n(\\epsilon)= 0.04$\\,eV\\,cm$^{-3}$ at $\\epsilon=0.1$\\,eV. \\\\ In the meantime, besides Mkn\\,421 (Punch et al. 1992, Petry et al. 1996) Mkn\\,501 (redshift $z$ = 0.034) is now the second extragalactic TeV source to be discovered and extensively monitored with the Whipple and HEGRA \\v{C}erenkov telescopes (Kerrick et al. 1995a, Bradbury et al. 1997). Both objects belong to the blazar subclass of galaxies showing powerful non-stellar activity characterized by rapidly variable, polarized continuum emission. As the outburst of Mkn\\,501 in early March 1997 (Aharonian et al. 1997) and the two outbursts of Mkn\\,421 in Spring of 1995 and 1996 (Kerrick et al. 1995b, Gaidos et al. 1996, Buckley et al. 1996) have shown, both objects also exhibit rapid variability with large amplitudes in the TeV energy range. The HEGRA collaboration showed that between March 16 and March 20 (about 27 hours observation time) the Mkn\\,501 energy spectrum in the TeV energy range extended beyond 10 TeV without a visible break in the power law spectrum with a differential power law index of 2.49 $\\pm$ 0.11 (stat.) $\\pm$ 0.25 (syst.) (Aharonian et al. 1997). During the quiescent state of Mkn\\,501 in 1996, this source was observed with the HEGRA CT1 telescope for 220 hours. The measured energy spectrum for this period can be described by a differential power law index of 2.5 $\\pm$ 0.4 (total error) (Petry 1997). \\\\ In the following we first discuss the current situation regarding models and observations of the DIRB and then derive an upper limit on the DIRB from the measured energy spectrum of Mkn\\,501 in 1997 for IR photon energies between 3$\\cdot 10^{-3}$ eV and 3$\\cdot 10^{-1}$ eV. ", + "conclusions": "Based on the observation of the {\\sl unabsorbed} Mkn\\,501 $\\gamma$-spectrum extending beyond 10\\,TeV, we have derived a stringent upper limit on the extragalactic diffuse infrared photon energy density in the energy range from $3 \\cdot 10^{-3}$ to $3 \\cdot 10^{-1}$\\,eV of 1.8 times the prediction of an average model by MacMinn \\& Primack. This translates into an upper limit of the energy density at the used pivot energy point of $3 \\cdot 10^{-2}$\\,eV of $\\epsilon^2 n(\\epsilon) = 1.1 \\cdot 10^{-3}$\\,eV/cm$^3$. For the second ansatz, i.\\,e.\\,, a power-law ansatz for the DIRB around this pivot energy point, we determined upper limits on the normalization as a function of the spectral index, e.\\,g.\\,for a flat spectrum $\\epsilon^2 n(\\epsilon) = n_0$, the resulting upper limit is $n_0 = 1.0 \\cdot 10^{-3}$\\,eV/cm$^3$. These upper limits are about 2 orders of magnitude {\\sl below} upper limits derived from current direct measurements at this energy and are not in contradiction to a value of the DIRB density derived from preliminary evidence of TeV emission reported by Meyer \\& Westerhoff ( Meyer \\& Westerhoff 1996). \\\\ The results presented in this paper are well compatible with other analyses of the infrared photon density based on the HEGRA CT data of Mkn\\,501 (Mannheim 1997, Stanev \\& Franceschini 1997). They are also in good agreement with an empirical calculation of the DIRB based on galaxy luminosity functions in the IR (Malkan \\& Stecker 1997) which in turn is in good agreement with recent DIRB models derived from star formation data (Guiderdoni et al. 1997). \\begin{ack} We thank K.\\,Mannheim, G\\\"ottingen, for useful discussions and R.\\,S.\\,Miller, Los Alamos, for helpful comments on the manuscript. \\end{ack}" + }, + "9802/astro-ph9802293_arXiv.txt": { + "abstract": "A tight mass-temperature relation, $M(r)/r\\propto T_X$, is expected in most cosmological models if clusters of galaxies are homologous and the intracluster gas is in global equilibrium with the dark matter. We here calibrate this relation using 8 clusters with well-defined global temperatures measured with {\\it ASCA} and masses inferred from weak and strong gravitational lensing. The surface lensing masses are deprojected in accordance with N-body simulations and analytic results. The data are well-fit by the mass-temperature relation and are consistent with the empirical normalisation found by Evrard \\etal\\ (1996) using gasdynamic simulations. Thus, there is no discrepancy between lensing and X-ray derived masses using this approach. The dispersion around the relation is 27 per cent, entirely dominated by observational errors. The next generation of X-ray telescopes combined with wide-field {\\it HST} imaging could provide a sensitive test of the normalisation and intrinsic scatter of the relation resulting in a powerful and expedient way of measuring masses of clusters of galaxies. In addition, as $M(r)/r$ (as derived from lensing) is dependent on the cosmological model at high redshift, the relation represents a new tool for determination of cosmological parameters, notably the cosmological constant $\\Lambda$. ", + "introduction": "Clusters of galaxies are the largest gravitationally bound structures in the Universe and are as such excellent probes of cosmic structure formation and evolution. The ensemble properties of clusters expected in various cosmological scenarios can be used to derive constraints on the power spectrum of the initial density perturbations and on cosmological parameters such as $\\Omega_0$ and $\\Lambda$ (e.g., Eke, Cole \\& Frenk 1996; Bahcall, Fan \\& Cen 1997; Oukbir \\& Blanchard 1997; Bartelmann \\etal\\ 1998; de Theije, van Kampen \\& Slijk\\-huis 1998). On the scales of individual clusters the inferred baryon mass fraction can be used to constrain $\\Omega_0$ (White \\etal\\ 1993; Evrard 1997). In such studies, an important quantity is the total cluster mass or any observed quantity which is tightly related to the mass. A promising mass estimator is the mean emission-weighted temperature, $T_X$, of the hot intracluster medium (ICM) in clusters of galaxies. Based on numerical simulations, it has been shown that $T_X$ is a better indicator of the total mass of a cluster than any other optical or X--ray property (Evrard 1990). Recently, Evrard, Metzler \\& Navarro (1996, hereafter EMN) and Eke, Navarro \\& Frenk (1997) showed that there is a tight relation between the mass of a cluster and its global X-ray temperature in cosmological gasdynamic simulations, irrespective of the state of the cluster (e.g., not restricted to clusters with a `regular' appearance or `isothermal' clusters) and the assumed cosmological model. In the simulations, it was found that mass predictions using this method (which only involve temperatures) are twice as precise as those derived using the $\\beta$-model (which require the surface brightness distribution in addition, i.e., more photons and higher spatial resolution). However, the normalisation of the relation hinges on numerical simulations which may not comprise sufficient detail (EMN; Anninos \\& Norman 1996). Therefore, it is essential to calibrate this relation from an observational point of view, by using independent mass estimators. The purpose of this \\paper\\ is to provide a first observational calibration of the $M$--$T_X$ relation using the relatively `clean' way of determining independent cluster masses by gravitational lensing. This technique essentially probes the projected mass along the line of sight. It is also pointed out that the relation holds the promise of providing a test of the geometry of the Universe which is particularly sensitive to $\\Lambda$. Throughout this \\paper, however, we assume a standard homogeneous Einstein--de Sitter Universe with $H_0=100 h$ km s$^{-1}$ Mpc$^{-1}$, $h=0.5$, $\\Omega_0=1$ and $\\Lambda = 0$. ", + "conclusions": "Based on numerical simulations (EMN; Eke \\etal\\ 1997) and observations of nearby clusters (Mohr \\& Evrard 1997) the existence of a tight mass-temperature relation has been suggested. The results presented here provide support for this assertion and indicate that the mass-temperature relation (eq.~6) can be used to determine cluster masses with a precision of 27 per cent (Fig.~2). There seems to be no significant discrepancy between deprojected lensing masses and masses derived from X-ray temperatures, using the normalisation found in numerical simulations (EMN). The origin of this tight relation is believed to be the fairly simple physics entering the relation (cf.~Sec.~2), namely virialisation of gravitationally bound structures with self-similar dark-matter density distributions that are in global quasi-equilibrium with the hot ICM, independent of the chosen world model, power spectrum, or exact formation redshift of the cluster. We have cautioned that the observational data discussed in this Letter are quite uncertain and possibly affected by systematic errors. The results should therefore only be taken as an indication of a tight mass-temperature relation. However, the future observational situation is promising. A sample of clusters with very precise lensing masses (e.g., from wide-field {\\it HST} imaging with the ACS) to about 10 per cent or better (e.g., Natarajan \\etal\\ 1998; Hoekstra \\etal\\ 1998) and equally accurate temperatures (e.g., with {\\it AXAF}, {\\it Spectrum--XG}, or {\\it XMM}) would allow us to study the intrinsic scatter of the relation and determine a precise normalisation. This could provide a direct and reliable mass estimator for distant clusters with important cosmological implications." + }, + "9802/astro-ph9802220_arXiv.txt": { + "abstract": "The Smith high velocity cloud (\\Vlsr\\ $=$ 98\\kms) has been observed at two locations in the emission lines \\OIIIr, \\NIIb\\ and \\Ha. Both the \\NII\\ and \\Ha\\ profiles show bright cores due to the Reynolds layer, and red wings with emission extending to \\Vlsr\\ $\\approx$ 130\\kms. This is the first simultaneous detection of two emission lines towards a high velocity cloud, allowing us to form the ratio of these line profiles as a function of LSR velocity. At both cloud positions, we see a clear distinction between emission at the cloud velocity, and the Reynolds layer emission (\\Vlsr\\ $\\approx$ 0). The \\NII/\\Ha\\ ratio ($\\approx$0.25) for the Reynolds layer is typical of the warm ionised medium. At the cloud velocity, this ratio is enhanced by a factor of $3-4$ compared to emission at rest with respect to the LSR. A moderately deep upper limit at \\OIII\\ (0.12R at 3$\\sigma$) was derived from our data. If the emission arises from dilute photoionisation from hot young stars, the highly enhanced \\NII/\\Ha\\ ratio, the \\OIII\\ non-detection and weak \\Ha\\ emission (0.24$-$0.30R) suggest that the Smith Cloud is 26$\\pm 4$ kpc from the Sun, at a Galactocentric radius of 20$\\pm 4$ kpc. This value assumes that the emission arises from an optically thick slab, with a covering fraction of unity as seen by the ionising photons, whose orientation is either (a) parallel to the Galactic disk, or (b) such as to maximize the received flux from the disk. The estimated mass and size of the cloud are $4\\times 10^6$\\Msun\\ and 6 kpc. We discuss a possible association with the much larger Sgr dwarf, at a galactocentric radius of 16$\\pm$2 kpc, which lies within 35$^\\circ$ ($\\sim$12 kpc) of the Smith Cloud. ", + "introduction": "\\subsection{Historical context} Since their discovery thirty years ago (Muller\\etal\\ 1963), the nature and origin of high velocity clouds (HVCs) has remained highly controversial. Their chequered history has been discussed by Verschuur (1988) and Wakker \\& van Woerden (1997). HVCs $-$ which cover at least a third of the sky $-$ are concentrations of neutral hydrogen with velocities which do not conform to a simple model of galactic rotation. Few, if any, clouds have reliable distance determinations, which has encouraged wide ranging speculation as to their origin. Explanations (assumed distances are given in parentheses) range from local supernova remnants ($\\sim$100 pc), large-scale expanding motions in nearby spiral arms ($<$ 1 kpc), condensations in the local galactic halo ($\\sim$ 1 kpc), structures in the galactic warp (5$-$20 kpc), tidal disruptions of the Magellanic Clouds (20$-$50 kpc), intergalactic gas ($>$ 50 kpc) or protogalaxies ($\\sim$500 kpc). The distance uncertainty continues to be the major stumbling block in understanding HVCs (Schwarz, Wakker \\& van Woerden 1995) since the cloud density and mass scale inversely with distance, and as the square of the distance, respectively. Ferrara \\& Field (1994; see also Wolfire\\etal\\ 1995) have suggested one possible method for systematic distance determinations based on the core/envelope structure observed in some HVCs (Cram \\& Giovanelli 1976). Part of the problem is that, besides \\HI\\ observations, it has proved difficult to detect HVCs in other spectral windows (Wakker \\& Boulanger 1986; Colgan, Salpeter \\& Terzian 1990). \\subsection{Previous experiments} Initially, the prospect of optical emission line detections of HVCs looked bleak. Reynolds (1987) did not detect any of six HVCs in the range \\Em\\ $=0.6-1$ cm$^{-6}$ pc with the Fabry-Perot `staring' technique. But there have since been many attempts to detect HVCs at wavelengths other than \\HI, with sensitive \\Ha\\ observations having the highest success rate. Wakker \\& van Woerden (1997) list the following published \\Ha\\ detections of HVCs: HVC 168-43-280 (0.08 R, Kutyrev \\& Reynolds 1989), cloud M~II (0.1-0.2 R; Munch \\& Pitz 1990), complex C (0.03 R, Songaila\\etal\\ 1989; 0.09 R, Tufte\\etal\\ 1998) and the Magellanic Stream (MS~II, MS~III, and MS~IV, at 0.37, 0.21, and 0.20 R respectively; Weiner \\& Williams 1996). Other optical detections of HVCs rely primarily on absorption line studies. This work can be used to constrain distances and determine metallicities (i.e. Magellanic metallicities in Lu\\etal\\ 1998; see also Schwarz, Wakker \\& van Woerden 1995). Most detections of visual band and/or ultraviolet absorption by metal ions are listed in Table 3 of Wakker \\& van Woerden (1997). The ions most commonly detected are \\CaII, \\MgII, and \\CIV. van Woerden\\etal\\ (1997) recently used Ca K absorption to determine an upper limit on the distance to complex A, giving 4$-$11 kpc as the first distance bracket for an HVC. HST is frequently used to search for the UV absorption lines, with the most recent detection by Sahu \\& Blades (1997) of \\SiII\\ towards HVC 487. All CO observations of HVCs have resulted in non-detections (Hulsbosch 1978; Giovanelli 1986; Kim\\etal\\ 1989), indicating that some HVCs are further than 3 kpc away. Though previous searches for far infrared emission have been negative (Wakker \\& Boulanger 1986; Bates\\etal\\ 1988; Fong\\etal\\ 1987), Ivesic \\& Christodoulou (1997) searched the IRAS point source catalog and found possible young stellar objects in the \\HI\\ cores of the M cloud and complex H. There have been claims for enhanced x-ray emission towards HVCs from ROSAT observations, \\eg towards Complex M near M~I and M~II (Herbstmeier 1995), and possibly from Complex C (Kerp 1996). Blom (1997) claims extended MeV emission is associated with HVC Complexes M and A, in the same location as associated diffuse soft x-rays. Absorption at 21 cm has been more successful for HVC detection. It yields a measurement of the hydrogen spin temperature, T$_s$, and can be used to determine the kinetic gas temperature (Dickey 1979; Liszt 1983). The results of this work generally yield temperatures between 50$-$100K. Colgan\\etal\\ (1990) finds T$_s$ $>$ 20-70K for the Anticentre Clouds, T$_s$ $>$ 20K for HVC 43-13-309, and T$_s$ $>$ 10K for the Magellanic Stream. Mebold (1991) find T$_s$ $>$ 50K for another position in the Magellanic Stream. Definitive results include two components at 70K and 300K in Cloud R (Payne\\etal\\ 1978; Payne\\etal\\ 1980) and 50K for complex H (Wakker\\etal\\ 1991). We now present the first simultaneous detection of more than one emission line towards an HVC. The observations presented here were first reported by Bland-Hawthorn (1994). In $\\S$2, we describe the observations before summarising the calibrations and reductions in $\\S$3. In $\\S$4, the high velocity cloud measurements are presented, and these are discussed in the context of ionisation models in $\\S$5. ", + "conclusions": "The recent detections of the Magellanic Stream in \\Ha\\ (Weiner \\& Williams 1996) provided Bland-Hawthorn \\& Maloney (1998) with the critical normalisation for the emergent UV flux from the Galactic disk. This model has been used to predict the ionising field throughout the Galactic halo. We determine a distance to the Smith Cloud of 26$\\pm$4 kpc on the basis of the \\Ha\\ flux alone. In support of this picture, the \\NII/\\Ha\\ is greatly enhanced at the velocity of the cloud and \\OIII\\ must be very weak, which together indicate a dilute ionising field impinging the cloud surface. Possible complications are cloud geometry, porosity, and uncertain extinction corrections. A lower mean disk opacity would put the clouds further away. We anticipate that our model is more reliable for clouds at greater vertical distances as this tends to average out structure in the distribution of UV sources. Clouds within a few kpc of the Galactic plane could be in relative shadow, particularly for opaque disk models. More accurate distances for HVCs will come from large \\HI\\ surveys with optical follow-up. A crucial development has been the availability of target lists with higher sensitivity and better resolution, in particular, the Parkes \\HI\\ Multi-beam Survey (Staveley-Smith 1997). More recent detections at the AAT are sufficiently strong to suggest that a significant number of clouds should be observable with this technique. As such, it provides a crucial test of the Blitz-Spergel model (Blitz\\etal\\ 1996; Spergel\\etal\\ 1996) which places roughly half of all high velocity clouds at extragalactic distances ($\\sim 1$ Mpc) within the Local Group, in which case none of the clouds should be detectable at these levels. We have set out to present a simplified picture as a challenge to theorists and experimentalists alike. The observational programmes that we have described here only require small to medium-sized telescopes (0.5$-$4m) since both HVC and Stream clouds subtend large angles, and the line surface brightness is expected to be fairly constant over large angular scales ($\\leq 1^\\circ$). Since the Fabry-Perot interferogram is binned azimuthally in order to produce the detection, the site does not require good seeing although it should have relatively good photometric stability. A dry site is favoured because variable water vapour features can complicate sky subtraction for \\Em $<$ 1 cm$^{-6}$ pc. We encourage a more widespread interest in the Fabry-Perot `staring' technique as it is set to make a profound contribution to the understanding of both galactic and extragalactic radiation fields over the coming years." + }, + "9802/astro-ph9802016_arXiv.txt": { + "abstract": "We present a new investigation of the Cloverleaf (z$=$2.56) based on the combination of archival HST/WFPC2 data, recent IRAM CO(7-6) maps and wide field CFHT/FOCAM images. The deep WFPC2 observation (F814W) shows a significant overdensity of I$_{814W}\\sim$ 23--25 galaxies around the Cloverleaf that we interpret as the presence of a distant cluster of galaxies along the line of sight. The Cloverleaf is probably the result of the lensing effects of a system which includes, in addition to a single galaxy, one of the most distant clusters of galaxies ever detected. We have modelled the lens accordingly. \\\\ The high-resolution IRAM/CO map provides for the first time the orientation and the ellipticity of the CO spots induced by the shear component. Velocity - positional effects are detected at the 8$\\sigma$ level in the CO map. A strong limit can then be put on the size, shape and location of the CO source around the quasar. The CO source is found to form a disk- or ring-like structure orbiting the central engine at $\\sim$ 100km/s at a radial distance of $\\sim$ 100pc, leading to a central mass of $\\sim$ 10$^{9}$ M$_\\odot$ possibly in the form of a black hole. ", + "introduction": "The Cloverleaf is the gravitationally lensed image of the quasar H1413+117 (14$^{\\rm h}$ 15$^{\\rm m}$ 46$^{\\rm s}$.23; $11^{\\rm o}\\ 29'\\ 44''.0$ J2000.0) at $z=2.558$ showing four spots with angular separations from $0''.77$ to $1''.36$. Since its discovery (Magain et al. 1988), the Cloverleaf has been observed spectroscopically and imaged with ground based telescopes in various bands from B to I as well as at 3.6 cm with the VLA. An early model of the gravitational lens was derived by Kayser et al (1990). The main difficulty with this model lies in the fact that it predicts, for a z=1.44 lens, a mass of $\\sim$ 2.5 $10^{11}$ $h^{-1}_{50}$ M$_\\odot$ within $0.7''$ (6 $h^{-1}_{50}$ kpc) radius, which would correspond to a relatively bright normal galaxy: so far, searches in the K band of the predicted `bright' lensing galaxy have been unsuccessful (Lawrence 1996) and this fact remains a mystery.\\\\ New data sets are available : the post-COSTAR HST/WFPC2 UV images (partly discussed in Turnshek et al. 1997), maps and spectral informations in the molecular CO transitions (Barvainis et al 1994, Wilner et al 1995, Barvainis et al 1997, Yun et al 1997, Alloin et al 1997) and CCD images of the Cloverleaf over a 5' field of view (CFHT/FOCAM archive) allowing a precise astrometry of the field obtained through different wavebands. New results obtained from these data are discussed in the following sections. Throughout the paper, we use H$_0$= 50 $h_{50}$ km/s/Mpc, $\\Omega_0$=1 and $\\Lambda=0$. \\begin{figure} \\vspace{1.75in} \\caption{Image of the Cloverleaf obtained with the IRAM telescope at Plateau de Bure. (a) is the total CLEANed image, (c) the CLEANed blue-shifted image, (d) the CLEANed red-shifted image and (b) the difference between the CLEANed red and CLEANed blue image. The CLEANed CO(7-6) maps were obtained with a natural beam of $0.8''\\times 0.4''$ at P.A. 15 deg. They have been restored with a circular $0.5''$ beam for comparison with HST data. Contour spacing is 1.35 mJy/beam for 1a, 2 mJy/beam for 1c and 1d and 3 mJy/beam for 1b, corresponding in each case to 2$\\sigma$. } \\label{fig:iram} \\end{figure} ", + "conclusions": "This new analysis of the Cloverleaf reveals that this is probably a complex lens which includes a lensing galaxy and an additional distant lensing cluster of galaxies. The reality of the cluster toward the Cloverleaf has still to be confirmed independently. Yet, most of the faint galaxies around the Cloverleaf are found in the same magnitude and size ranges, as expected if they indeed belonged to a cluster. If the cluster is at a very large distance, the shift of its galaxy luminosity function up to higher apparent magnitude would explain why the number-density contrast of the cluster with respect to faint field galaxies is lowered down to only a 4 $\\sigma$ level. \\\\ This interpetation implies that the lensing galaxy may not be very massive and consequently may not be very luminous, helping to explain the mystery of the lensing galaxy not having been detected so far. The drawback is that, despite the constraint that the shapes of the CO spots provide on the orientation of the mass density distribution, it mandates sharing the mass between the lensing-galaxy and the lensing-cluster which increases the number of possible lens configurations. Further data might reveal the position and the shape of the light distribution of the lensing galaxy.\\\\ Such information would be useful to improve the mapping of the CO source, as emphasized by Alloin et al. (1997). With the present-day data, the CO source is found to be a disk- or ring-like structure with typical radius of $\\sim$ 100 pc, under the lens composite model of a galaxy and a cluster at z=1.7, leading to a central $\\sim$ 10$^{9}$ M$_\\odot$ object, typical of a massive black-hole. It is amazing to see that a disk with such a small intrinsic size can be spatially ``resolved\" even at an angular distance as large as 1.6 $h^{-1}_{50}$ Gpc.\\\\ Although this remains to be confirmed independently, the discovery of a distant cluster of galaxies on the line of sight to the Cloverleaf is remarkable because it reinforces the suspicion that many bright high redshift quasars are magnified by cluster-like systems at large distances. This was already reported from analyses in the fields of the doubly imaged quasar Q2345+007 (Bonnet et al 1993; Mellier et al 1994; van Waerbeke et al 1997), where a cluster candidate is expected to be at z$\\sim$ 0.75 (Pell\\'o et al 1996) and of MG2016 where the X-ray emission of the intra-cluster gas has been observed and for which the Iron line (from X-ray spectroscopy) gives a redshift z$\\sim$1 (Hattori et al 1997)." + }, + "9802/astro-ph9802199_arXiv.txt": { + "abstract": "I have undertaken a literature search through 31 July 1997 of white dwarfs (WDs) in open and globular clusters. I have tried to make a careful evaluation in each case of the likelihood that the object is a white dwarf and that it is a cluster member. The results are presented for 13 open clusters and 11 globular clusters. Currently there are 36 single WDs and 5 WDs in binaries known among the open clusters, and 340 single WDs and 11 WDs in binaries known among the globular clusters. From these data I have calculated WD mass fractions for four open clusters (the Pleiades, NGC 2168, NGC 3532, and the Hyades) and one globular cluster (NGC 6121). I develop a simple model of cluster evolution that incorporates stellar evolution but not dynamical evolution to interpret the WD mass fractions. I augment the results of my simple model by turning to sophisticated N-body simulations incorporating stellar evolution (Terlevich 1987; de la Feunte Marcos 1996; Vesperini \\& Heggie 1997). I find that even though these clusters undergo a range of degrees of kinematical evolution from moderate (the Pleiades, NGC 2168, and NGC 3532) to strong (the Hyades and NGC 6121) the WD mass fraction is relatively insensitive to kinematical evolution and little changed from a model incorporating only stellar evolution with a Salpeter-like initial mass function. By comparing the cluster mass functions to that of the Galactic disk, and incorporating plausibility arguments for the mass function of the Galactic halo, I estimate the WD mass fraction in these two field populations. I assume the Galactic disk is $\\sim10$ Gyrs old (Winget \\etal 1987; Liebert, Dahn, \\& Monet 1988; Oswalt \\etal 1996) and that the Galactic halo is $\\sim12$ Gyrs old (Reid 1997b; Gratton \\etal 1997; Chaboyer \\etal 1998), although the WD mass fraction is insensitive to age within this regime. I find that the Galactic halo should contain from $8$ to $9$\\% ($\\alpha = -2.35$) or perhaps as much as $15$ to $17$\\% ($\\alpha = -2.0$) of its stellar mass in the form of WDs. The Galactic disk WD mass fraction should be $6$ to $7$\\% (for a median stellar age of $5$ to $7$ Gyrs and $\\alpha = -2.35$), consistent with the empirical estimates of $3$ to $7$\\% (Liebert, Dahn, \\& Monet 1988; Oswalt \\etal 1996). ", + "introduction": "Since white dwarfs are faint for most of their evolutionary lifetime, their mass fraction in clusters and in the field is difficult to measure. Yet the WD mass fraction is important both for the dynamical evolution of star clusters and potentially for the mass of the Galactic disk and halo. Even in the immediate solar neighborhood, the range of the WD mass density estimates vary by more than a factor of two, from $2.0 \\times 10^{-3} M_{\\sun} \\ pc^{-3}$ (Liebert, Dahn, \\& Monet 1988) to $4.6^{+2.2}_{-0.4} \\times 10^{-3} M_{\\sun} \\ pc^{-3}$ (Oswalt \\etal 1996). While the solar neighborhood stellar density itself is poorly constrained, for a value of $\\sim6.4 \\times 10^{-2} M_{\\sun} \\ pc^{-3}$ (Mihalas \\& Binney 1981; and consistent with Kuijken \\& Gilmore 1989, after subtracting the interstellar gas mass) the WD mass fraction ranges from $3$ to $7$\\%. In the Galactic halo the situation is even more poorly constrained, and the WD mass fraction is effectively observationally unknown. Indeed, studies of gravitational lensing in the Milky Way (e.g.\\ Alcock \\etal 1997) led to a flurry of papers during 1997 examining whether $\\sim50$\\% of the Galactic dark matter could be in the form of halo WDs. The bulk of these studies concluded that such a high halo WD mass fraction can be ruled out (see Gibson \\& Mould 1997, and references therein), but the mere fact that the mass fraction of WDs is so poorly known drives speculation about its importance. For the clusters, the presumed source of the field WDs, the WD mass fraction must depend on the cluster age and kinematical evolution (e.g.\\ Vesperini \\& Heggie 1997). In the last three years a number of studies have identified and measured the properties of WDs in open and globular clusters. Most of these new cluster WD measurements have been made possible by the ability of the Hubble Space Telescope (HST) to detect very faint point sources and separate them from the many faint resolved background galaxies. These studies have been motivated by the independent information available from cluster WDs on cluster distances (Renzini \\etal 1996), cluster ages (von Hippel, Gilmore, \\& Jones 1995), and constraints on stellar evolution (Richer \\etal 1997). An important byproduct of these studies is the number and mass contribution of WDs to their parent clusters. To the best of my knowledge no one has yet extracted this important information. In this paper I first tabulate the known cluster WDs and estimate their fraction by mass in a handful of clusters. I then use a simple interpretive model supplemented by cluster dynamical studies in the literature to argue that the observed numbers of cluster WDs are about what one would expect based on stellar evolution theory alone and are insensitive to the cluster dynamical history. Finally, I discuss the relevance of the cluster WD mass fractions to the disk and halo field star WD mass fractions. ", + "conclusions": "I have undertaken a literature search through 31 July 1997 of white dwarfs (WDs) in open and globular clusters. I have tried to make a careful evaluation in each case of the likelihood that the object is a white dwarf and that it is a cluster member. The results are presented for 13 open clusters and 11 globular clusters. Currently there are 36 single WDs and 5 WDs in binaries known among the open clusters, and 340 single WDs and 11 WDs in binaries known among the globular clusters. From these data I have calculated WD mass fractions for four open clusters (the Pleiades, NGC 2168, NGC 3532, and the Hyades) and one globular cluster (NGC 6121). I develop a simple model of cluster evolution that incorporates stellar evolution but not dynamical evolution to interpret the WD mass fractions. I augment the results of my simple model by turning to sophisticated N-body simulations incorporating stellar evolution (Terlevich 1987; de la Feunte Marcos 1996; Vesperini \\& Heggie 1997). I find that even though these clusters undergo a range of degrees of kinematical evolution from moderate (the Pleiades, NGC 2168, and NGC 3532) to strong (the Hyades, NGC 6121) the WD mass fraction is relatively insensitive to kinematical evolution and little changed from a model incorporating only stellar evolution with a Salpeter-like initial mass function. By comparing the cluster mass functions to that of the Galactic disk, and incorporating plausibility arguments for the mass function of the Galactic halo, I estimate the WD mass fraction in these two field populations. I assume the Galactic disk is $\\sim10$ Gyrs old (Winget \\etal 1987; Liebert, Dahn, \\& Monet 1988; Oswalt \\etal 1996) and that the Galactic halo is $\\sim12$ Gyrs old (Reid 1997b; Gratton \\etal 1997; Chaboyer \\etal 1998), although the WD mass fraction is insensitive to age within this regime. I find that the Galactic halo should contain from $8$ to $9$\\% ($\\alpha = -2.35$) or perhaps as much as $15$ to $17$\\% ($\\alpha = -2.0$) of its stellar mass in the form of WDs. The Galactic disk WD mass fraction should be $6$ to $7$\\% (for a median stellar age of $5$ to $7$ Gyrs and $\\alpha = -2.35$), consistent with the empirical estimates of $3$ to $7$\\% (Liebert, Dahn, \\& Monet 1988; Oswalt \\etal 1996). Ultimately, precise comparisons between the field and cluster MFs for both the disk and halo would be a means of determining if the clusters we see today are typical of those that built the field populations." + }, + "9802/astro-ph9802150_arXiv.txt": { + "abstract": "We discuss the spectral character of Lyman--limit--selected, star--forming galaxies at $z > 3$. The rest--frame UV spectra of these faint galaxies may show Ly$\\alpha$ in either absorption or emission, probably depending upon their local ISM content and geometry. Other UV interstellar resonance absorption lines show considerable variation in strength, likely related to differences in the galactic metal abundances. We present initial results on $B$--drop galaxies, generally at $z\\sim 4$. Our low--resolution spectrograms show no measurable flux below the redshifted Lyman limit (912\\,\\AA). Thus, it is likely that normal, star--forming galaxies at early cosmic epochs did not significantly contribute to the metagalactic ionizing radiation field. ", + "introduction": "While not the ``main characters'' in the Academy colloquium, actively star--forming galaxies are now routinely discovered and made available for study through color--selection. $U$-- and $B$--dropouts, galaxies targeted for the redshifted Lyman limit spectral discontinuity at 912 \\AA\\ in the $U$-- and $B$--bands, are important in the early Universe because of their ubiquity ({\\em c.f.\\ }Steidel {\\em et al.\\ }1996a,b). The $U$--drops are several percent of the total deep number counts at $R \\sim 24^{m}$ (corresponding to $B \\sim 25^{m}$), with an absolute surface density of $\\sim 3$ galaxies per square arcminute on the sky. Moderate power radio sources are far less numerous, with surface densities of $\\sim 2\\times 10^{-3}$ radio sources per square arcminute at $S_{\\rm 1.4~GHz} = 10$\\,mJy. In an informal collaboration with C.~Steidel and M.~Pettini, we have pursued moderate resolution observations of the brightest ($R \\sim 23^{m}$) $z\\sim 3$ star--forming galaxies discovered by Steidel and collaborators ({\\em e.g.}, Steidel \\& Hamilton 1992; Steidel {\\em et al.\\ }1996a). Typical sources are at $z\\sim 3$ and require long integrations at Keck with the LRIS spectrograph. For the fainter and (usually) more distant $B$--drop galaxies ($z\\sim 4$), lower spectral resolution is obtained. This work is part of an ongoing effort to study `normal', {\\em i.e.}, non--AGN, young galaxies at the earliest cosmic epochs. \\begin{figure}[t] \\epsfxsize=5in \\epsffile{spinradfig1.ps} \\caption{Sequence of UV spectra of four Lyman--limit galaxies, arranged from highest redshift to lowest redshift. Note the large range in line strengths present. In particular, Ly$\\alpha$ ($\\lambda\\,1216$\\,\\AA) is visible in both emission and absorption, while primarily interstellar resonant lines (see Table~1) show considerable variation in strength, possibly related to metallicity in these young star--forming galaxies. The Ly$\\alpha$ absorption galaxies also appear redder in their continua at $\\lambda>1300$\\,\\AA .} \\end{figure} ", + "conclusions": "" + }, + "9802/astro-ph9802201_arXiv.txt": { + "abstract": "The details of \\v Cerenkov light produced by a $\\gamma -$ ray or a cosmic ray incident at the top of the atmosphere is best studied through systematic simulations of the extensive air showers. Recently such studies have become all the more important in view of the various techniques resulting from such studies, to distinguish $\\gamma -$ ray initiated showers from those generated by much more abundant hadronic component of cosmic rays. We have carried out here such systematic simulation studies using CORSIKA package in order to understand the \\v Cerenkov photon density fluctuations for 5 different energies at various core distances both for $\\gamma -$ ray and proton primaries incident vertically at the top of the atmosphere. Such a systematic comparison of shower to shower density fluctuations for $\\gamma -$ ray and proton primaries is carried out for the first time here. It is found that the density fluctuations are significantly non-Poissonian. Such fluctuations are much more pronounced in the proton primaries than $\\gamma -$ ray primaries at all energies. The processes that contribute significantly to the observed fluctuations have been identified. It has been found that significant contribution to fluctuations comes from photons emitted after shower maximum. The electron number fluctuations and correlated emission of \\v Cerenkov photons are mainly responsible for the observed fluctuations. ", + "introduction": "Ground based atmospheric \\v Cerenkov technique is, at present, the only way by which TeV $\\gamma -$ rays could be detected from point sources such as $\\gamma -$ ray pulsars, short period X$-$ ray binaries or BL-Lac objects. Recent detection of TeV emission from a few of these objects (Vacanti {\\it et al.,} 1991; Punch {\\it et al.,} 1992, Chadwick {\\it et al., }1997, Quinn {\\it et al.,} 1997, Weekes, 1988, Fegan, 1994) has created much interest in the field of TeV $\\gamma -$ ray astronomy. The $% \\gamma -$ ray signals found typically are $\\sim 1\\%$ of the more abundant background events of cosmic ray nuclei, particularly protons. In order to detect faint Very High Energy (VHE) $\\gamma -$ ray sources, one has to improve the signal to noise ratio by rejecting the bulk of the hadronic background. In order to do so it is imperative to study the detailed characteristics of \\v Cerenkov light production by photon initiated and proton initiated cascades in the atmosphere. Simulation studies in the past (Rao \\& Sinha, 1988; Hillas \\& Patterson, 1987; Zatsepin \\& Chudakov, 1962) have shown that the \\v Cerenkov pool at the observation level has the signature of the primary. The lateral distribution of \\v Cerenkov radiation seems to be distinctly different in $% \\gamma -$ ray and proton initiated showers in the sense that in the former case it is flat upto about 140 m and characterized by an increased photon density (called the `hump') at that distance while in the latter case it is steeper with practically no hump. It has been suggested (Rao \\& Sinha, 1988) that this characteristic difference could be measured in an observation and could be used for improving the signal to noise ratio. These arguments are based on the average properties of showers. In practice, however, \\v Cerenkov photon density fluctuations play a dominant role. The signature of the primary gets buried in the noise which is mainly due to large fluctuations in photon densities from shower to shower, even at same energy. These fluctuations in turn reduce the efficiency by which the primary could be identified based on the lateral distribution measurements (Krys and Wasilewski, 1993). As a result, the study of the photon fluctuations plays an important role in deciding the signal to noise ratio of the VHE $\\gamma -$ ray observations based on the measurement of lateral distribution of \\v Cerenokov light. It has been known that shower to shower fluctuations in proton initiated cascades are expected to be much larger than $\\gamma -$ ray initiated cascades since the nuclear interaction mean free path (70 $g$ $cm^{-2}$) is about twice as large as the radiation length (37.15 $g$ $cm^{-2}$) in the atmosphere and the number of secondaries and their energy spectra are known to fluctuate widely. Furthermore muons, which are present only in hadron initiated showers (above the \\v Cerenkov threshold $E_\\mu \\sim 4$ GeV) reach the observation level and could create local peaks in the light pool at the observation level (Hillas \\& Patterson, 1987). In this paper we make an attempt to estimate and compare the extent of these fluctuations in proton and $\\gamma -$ ray initiated showers at various energies and estimate the relative contributions from different but known sources of fluctuations. ", + "conclusions": "A systematic study of the fluctuations in the \\v Cerenkov photons generated by $\\gamma -$ ray and proton primaries in the earth's atmosphere and detected at the observation level has been carried out. Such a quantitative estimate of the degree of fluctuations for the two types of primaries and the dependence on the core distance as well as primary energy has been done for the first time here. This type of study is important in planning observations of VHE $\\gamma -$ ray sources based on the measurement of lateral distributions of \\v Cerenkov photons, since these experiments are based on improving signal to noise ratio by rejecting the abundant charged particle background. The large non-statistical fluctuations reported here might reduce the efficiency of rejection. We have studied the density fluctuations as a function of core distance for various energies of $\\gamma -$ ray and proton primaries in 50--1000 GeV range. Fluctuations are highly non-statistical and decrease with increasing primary energy in both the cases. Proton primaries show larger fluctuations compared to $\\gamma -$ ray primaries of corresponding energy. In case of $\\gamma -$ ray primaries fluctuations are minimum at the hump region and approach Poissonian beyond the hump region. Proton showers in general show larger fluctuations than $\\gamma -$ ray primaries even in the case of shower size measured in terms of the total number of \\v Cerenkov photons generated. We have investigated various known sources of fluctuations and tried to evaluate their relative contributions. Effect of geomagnetic field is to deflect electrons and thereby increase the fluctuations. Average electron energies and their spectra at different atmospheric levels during the shower development are found to be similar in case of proton and $\\gamma -$ ray primaries of equivalent \\v Cerenkov yields. Whereas number of electrons at various depths of shower development is found to vary much more for protons compared to $\\gamma -$ rays. We have also studied the effect of variation in the first point of interaction on the shower size fluctuations. It is found to be significant only in the case of $\\gamma -$ ray primaries. Contribution to the fluctuations also comes from Coulomb scattering of low energy electrons past the shower maximum and the intrinsic correlation between the photons emitted by a single electron." + }, + "9802/astro-ph9802171_arXiv.txt": { + "abstract": " ", + "introduction": "The study of the cooling of neutron stars is being pursued principally with the hope to constrain the structure of dense matter since young neutron stars cool by neutrino emission which comes mainly from the matter at supranuclear densities in the deep interior. The initial scenario, known as `The Standard Model', was based on the modified Urca process. This inefficient process leads to a slow cooling and the lack of detection of neutron stars, and their surface thermal emission, in several supernova remnants during the '60 and '70 stimulated the search for stronger neutrino emission channels. In 1965 Bahcall \\& Wolf had shown that free pions, if present, would enormously increase the neutrino emissivity: a realistic mechanism for the presence of pions, pion condensation, was soon proposed and with it the dichotomy between `standard' and `exotic' cooling scenarios (pions have nothing exotic, of course, but matter with a pion condensate is exotic). With the years, other possible fast cooling mechanisms have been proposed: quark matter, kaon condensation, the direct Urca process with nucleons and finally the direct Urca process with hyperons (see Pethick 1992 for a review). Naively, the standard model, on one side, and the fast cooling scenarios, on the other, lead to very different predictions and there was great hope to be able to distinguish them by comparison of cooling models with data: fast cooling scenarios predict much lower surface temperatures at ages below $\\sim 10^{5 - 6}$ yrs. In this way, a cold (say $T_e << 10^6$ K) young (say $t < 10^5$ yrs) neutron star would require fast neutrino cooling while a warm one would be evidence in favor of the standard model. However, things are more complicated than just the raw neutrino emission rate. Baryon pairing can suppress the neutrino emission when the temperature $T$ is below the critical temperature $T_c$: a high $T_c$ can lead to such strong suppression that, even with the most efficient neutrino emission mechanism (the direct Urca process) the resulting surface temperature can be higher than predicted within the standard model (Page \\& Applegate 1992). Unfortunately, the uncertainty on the value of $T_c$ in the core is such that almost any surface temperature could be accommodated within almost any fast cooling scenario. With the launch of ROSAT came the confirmation of the detection of surface thermal emission from several neutron stars (\\\"Ogelman 1995): the estimated surface temperatures lay close to the predictions of the standard model and made some theorists (particularly myself) quite unhappy. Later, a new and arguably more realistic fit of the Vela pulsar spectrum nevertheless gave a temperature estimate twice lower than the original one, which seemed to seriously indicate the occurrence of some fast cooling agent at least in that neutron star (Page et al. 1996). But it turned out that an essential process had been completely forgotten in all cooling calculations: neutrino emission due to the formation of Cooper pairs when the neutron (or proton) liquid becomes superfluid (resp. superconducting). This process is perfectly `standard' and, once taken into account, the standard model can be perfectly compatible with our present interpretation of the data: {\\em the standard model is alive and well alive} (see Schaab et al 1997; Page 1998; Yakovlev {\\em et al.} 1998). Non standard scenarios are of course also compatible with the data. I have presented earlier (Page 1998) a less sketchy discussion of the above lines, to which I refer the reader for more details and references. In the next sections I consider some complementary issues. \\begin{figure} \\vspace{7.5cm} \\special{psfile=fig1.eps} \\caption{{\\bf Neutron stars as thermometers for nucleon pairing:} see text. In B, the numbers on the curves indicate the assumed uniform value of the neutron $^3$P$_2$ pairing $T_c$, in units of 10$^9$ K, and two standard cooling curves are also shown for comparison. (From Page 1995). \\label{fig1}} \\end{figure} ", + "conclusions": "" + }, + "9802/astro-ph9802347_arXiv.txt": { + "abstract": "We discuss BIMA (Berkeley Illinois Maryland Association) data and present new high quality optical and near-IR Keck images of the bright radio ring PKS~1830--211. Applying a powerful new deconvolution algorithm we have been able to identify both images of the radio source. In addition we recover an extended source in the optical, consistent with the expected location of the lensing galaxy. The source counterparts are very red, $I-K\\sim7$ suggesting strong Galactic absorption with additional absorption by the lensing galaxy at $z=0.885$, and consistent with the detection of high redshift molecules in the lens. ", + "introduction": "The bright radio source PKS~1830--211 (Jauncey et al. 1991) has attracted much attention as the most detailed example of a lensed radio ring. Of the classically-lensed QSOs, its short time delay of 44 days (van Ommen et al. 1995) and clean lens geometry (e.g. Subrahmanyan et al, 1990; hereafter S90), make it a good candidate for measuring H$_{0}$. The lens, a gas rich galaxy at z=0.89, was discovered in the millimeter via molecular absorption (Wiklind \\& Combes 1996), which is seen towards only one of the two flat spectrum hot spots (Wiklind \\& Combes 1996, Frye et al. 1997). A nearby saturated M-star and heavy extinction along the line of sight (b=-5.7 degrees, Djorgovski, et al. 1992; hereafter D92) has obscured the lens and the source from identification. In this paper we describe how the MCS deconvolution algorithm (Magain, Courbin \\& Sohy, 1997) was used to detect the counterparts of this bright radio ring lens in deep Keck optical and infrared images. ", + "conclusions": "The main result of our study is the detection of the optical and near-IR counterpart to the NE radio source of PKS~1830-211 and the possible detection of the SW component and lensing galaxy. However, our SW component candidate might be the lensing galaxy alone (its position is also in good agreement with the predictions from the models calculated by S90 and Nair, Narasimha \\& Rao, 1993) or, given the crowding in the field, a red galactic object almost coincident with the position of the SW radio source. The hypothesis of a demagnified third image of the source (S90) between the 2 main lensed images is unlikely as in such a case extinction of the lens would have made it visible in the IR. Furthermore, the IR centroid of the SW component would have been shifted towards ``object E'' rather than to the radio position of the SW component. The higher contrast between the NE component and nearby M-star in the infrared make near-IR spectroscopy necessary for finding the source z. Deep, high resolution near-IR imaging is needed to reveal the exact nature of the faint SW component. However, even at the highest resolution attainable, 0\\farcs2- 0\\farcs3 in the IR with HST (in particular in K where the SW component of PKS1830-211 is best visible), deconvolution will be essential to discriminate between the SW component candidate, the lensing galaxy and additional faint galactic stars." + }, + "9802/astro-ph9802353_arXiv.txt": { + "abstract": "An inversion technique has been developed to recover LTE one-dimensional model photospheres for late-type stars from very high resolution, high signal-to-noise stellar line profiles. It is successfully applied to the Sun by using a set of clean \\ion{Ti}{1}, \\ion{Ca}{1}, \\ion{Cr}{1}, and \\ion{Fe}{1} normalized line profiles with accurate transition probabilities and taking advantage of the well understood collisional enhancement of the wings of the \\ion{Ca}{1} line at 6162 \\AA. Line and continuum center-to-limb variations, continuum flux, and wings of strong metal lines are synthesized by means of the model obtained and are compared with solar observations, as well as with predictions from other well known theoretical and empirical solar models, showing the reliability of the inversion procedure. The prospects for and limitations of the application of this method to other late-type stars are discussed. ", + "introduction": "\\label{sec1} Model photospheres for late-type stars are a fundamental cornerstone of modern Astronomy. The Sun, our closest star, belongs to this stellar category and its atmosphere provides an exceptional plasma laboratory to study detailed physical processes. Late-type stars are also the most numerous group in the Galaxy. They not only dominate its present dynamics, but their photospheric abundances constitute a unique tool for tracing its chemical evolution history. Furthermore, the oldest unevolved stars in the Galaxy, which provide information about the physical conditions at early epochs and constraints on models of primordial nucleosynthesis, are also late-type stars. In this century, considerable efforts have been devoted to theoretical and empirical modeling of late-type stellar atmospheres. In this respect, particular success was achieved in the seventies, and most of the models used nowadays are mainly based on work carried out in that decade. For a comprehensive review of the state of the art in the modeling of photospheres of F- and later-type stars see, for example, the discussion by Gustafsson \\& J\\o rgensen (1994) and Avrett (1996). Classical theoretical models are usually computed under the simplifying assumptions of hydrostatic equilibrium, plane-parallel stratification, local thermodynamic equilibrium (LTE), and conservation of energy flux (treating convection with the mixing length approximation). High-resolution, high signal-to-noise ratio spectra over extensive spectral ranges can be used to test the validity of existing theoretical models and to better constrain the observational features which have to be reproduced (line profiles and asymmetries, excitation and ionization equilibria of different elements, stellar fluxes, etc.). These data can also be used to construct models by using observed features to find a temperature--depth ($T-\\tau$) relation. Semi-empirical models of this kind have been obtained for the solar photosphere, where the present-day knowledge of the abundances of different elements, as well as its extended disk, allow a better observational constraint. The model constructed by Holweger \\& M\\\"uller (1974) is a good example of this class and has been widely used not only for the Sun but also for other late-type stars by scaling its $T-\\tau$ relation to different stellar effective temperatures. Although the spatially resolved information that can be obtained for the Sun is lacking for other late-type stars, spectral lines do contain enough information to determine reliable values for the physical magnitudes governing the state of a stellar atmosphere, and this approach has been applied, for instance, to the modelling of giant stars such as Arcturus (M\\\"ackle et al. 1975) and Pollux (Ruland et al. 1980). The so-called spectroscopic {\\it inversion} methods are aimed at obtaining the model atmosphere from which the synthetic line profiles best match the observed ones (usually with a least-squares criterion). The availability of high-quality data (observed spectra and line parameters) and modern computers make possible these inversions by taking into account the information contained in the whole line profile, which can potentially provide not only the temperature stratification in the atmosphere but also insights into other physical magnitudes such as velocity and magnetic fields. This paper is part of a wider project aimed at better understanding the photospheres of late-type stars of different metallicities for use in studies of Galactic chemical evolution and primordial nucleosynthesis. Among the different approaches envisaged within this project, which incorporates the study of theoretical model atmospheres and the presence and extension of line asymmetries in the photospheres of metal-poor stars, we present here a successful attempt to obtain a semi-empirical solar model photosphere by inverting observed clean line profiles with accurate oscillator strengths. This kind of information is also available for other late-type stars (with different levels of metallicity), and what we learn from modeling the Sun will help us when trying to apply the method to other stars. Section 2 provides the details of the inversion code, which is based on a well-tested inversion method for the Stokes parameter vector; \\S 3 shows its application to the solar case and the comparison of its results with those associated with other well-known theoretical and semi-empirical models. We describe in \\S 4 the potential use and limitations in applying the code to other late-type stars, and the main conclusions and the perspectives for the near future are summarized in \\S 5. ", + "conclusions": "\\label{sec5} The Multiline Inversion of Stellar Spectra (MISS) procedure developed here has been proved able to find an LTE temperature stratification for the solar photosphere from normalized high-resolution line profiles which reproduces the solar continuum and the limb-darkening sufficiently well and is in excellent agreement with observed strong and weak spectral lines with accurate line data. A better comparison of its performance in reproducing the limb-darkening of the wings of strong lines with accurate damping constants is now being studied. It is found that none of the input abundances (Anders \\& Grevesse 1989) for chromium, calcium, and titanium need to be changed to find the best fit to the observations, but the iron abundance is preferred to vary from the initially assumed value (7.67, in the scale where the hydrogen abundance is 12) to the lower (approximately meteoritic) abundance 7.48. A better fit to the solar flux line spectrum would require the abandonment of the hypothesis of hydrostatic equilibrium and the introduction of multi-component models or velocity fields inducing asymmetries in the line profiles. Our work in the near future will consider this improvement. This procedure of semi-empirical modeling of stellar photospheres can potentially be applied to other late-type stars. While other classical methods of empirical modeling would require observations not available for stars (the limb-darkening cannot be measured yet, and flux calibrated spectra are affected by a number of larger uncertainties), the normalized line profiles used by MISS are already at hand." + }, + "9802/astro-ph9802215_arXiv.txt": { + "abstract": "There is both theoretical and empirical evidence that the initial mass function (IMF) may be a function of the local star formation conditions. In particular, the IMF is predicted to flatten with increasing local luminosity density $\\rho _{l}$, with the formation of massive stars being preferentially enhanced in brighter regions. In R136, the bright stellar cluster in 30 Doradus, the IMF gradient is $\\partial \\Gamma /\\partial \\log \\rho _{{\\scriptsize\\rm l}}=0.28\\pm 0.06,$ where $\\Gamma $ is the slope of the IMF. If such IMF gradients are indeed general features of galaxies, this implies that several previous astrophysical measurements, such as the surface mass densities of spirals (obtained assuming constant mass to light ratios), were plagued by substantial systematic errors. In this Letter, calculations which account for possible IMF gradients are presented of surface densities of spiral galaxies. Compared to previous estimates, the mass surface densities corrected for IMF gradients are higher in the outer regions of the disks. For a model based on the Milky Way but with an IMF scaled according to R136, the rotation curve without the traditional dark halo component falls with Galactocentric radius, though slower than it would without IMF gradients. For a second model of the Milky Way in which the IMF gradient is increased to 0.42, the rotation curve is approximately flat in the outer disk, with a rotational velocity below $\\simeq 220$ km s$^{{\\scriptsize\\rm -1}}$ only before the traditional dark halo component is added. For a third model in which substantial arm/interarm density contrasts are additionally assumed, the solar vicinity mass density drops to $0.10 M_{\\odot }$pc$^{-3}$, which is consistent with observations. These results, if generalizable to other galaxies, not only call into question the assertion that dark matter halos are compatible with the flat rotation curves of spiral galaxies, but also may clarify our understanding of a wide variety of other astrophysical phenomena such as the G-dwarf problem, metallicity gradients, and the Tully-Fisher relation. ", + "introduction": "\\def\\figdirprefix{/home/theory1/taylor/stellarmodels/} In a recent paper, Padoan, Nordlund, \\& Jones (1997) claimed on theoretical grounds that the initial mass function (IMF) should be a function of the local temperature ${\\it T}$ of the original molecular clouds. Padoan et al. (1997) argued that dense star forming regions, such as those in starburst galaxies, should be warmer than sparser star forming regions. In fact, if the temperature dependence of the clouds is not drastically different from that of a blackbody, then $T\\propto \\rho ^{1/4}_{{\\scriptsize\\rm l}}$, where $\\rho _{{\\scriptsize\\rm l}}$ is the local mean luminosity once star formation has already started. Padoan et al. claimed that starburst regions should therefore have a flatter IMF and be more ``top heavy.\" Similar reasoning would imply that the IMF in cooler regions of galaxies should favor low mass star formation and be steeper. \\par In support of their star formation model, Padoan et al. (1997) noted that for $T\\gtrsim 60$ K, their models predict a top heavy IMF similar to that found in the center of R136~(Malumuth \\& Heap 1994, Brandl et al. 1996), the bright stellar cluster in 30 Doradus. Due to its proximity and the fact that it is the most massive H {\\scriptsize\\rm II} region in the Local Group, 30 Doradus is perhaps the best star formation ``laboratory\" accessible to us. However, the relaxation time in R136 may be less than its age (Campbell et al. 1992), so dynamic friction may also contribute toward the R136 present-day mass function gradient. \\par Fortunately, many other avenues of testing Padoan et al.'s model exist. The O-star catalog of Garmany, Conti, \\& Chiosi (1982) shows a flattening of the IMF slope toward the Galactic center (cf Humphreys \\& McElroy 1984). This data supports Padoan et al.'s model since higher surface brightness regions would, on the average, yield higher temperatures and flatter IMFs. Models which attempt to explain correlations between local surface brightness, color, line ratios, metallicity, and the star formation rate have assumed luminosity-dependent IMFs (e.g., Edmunds \\& Phillipps 1989; Phillipps, Edmunds, \\& Davies 1990). Several evolutionary models of inner regions of starburst galaxies assume low mass cutoffs or top heavy IMFs (e.g., Rieke et al. 1980; Augarde \\& Lequeux 1985; Doane \\& Mathews 1993; Doyon, Joseph, \\& Wright 1994). Finally, independent theoretical arguments supporting IMF gradients range from models which are consistent with the simple form of the Jeans expression for the typical stellar mass in solar units of $\\propto T^{3/2}~($e.g.; Larson 1982; Bodenheimer, Tohline, \\& Black 1980) to much more complicated models, such as the outflow-regulated model of Adams \\& Fatuzzo (1996), which predicts $\\propto T^{a}$, where $1\\le a\\le 3/2.$ \\par If IMFs are actually a function of $\\rho _{{\\scriptsize\\rm l}}$ or $T$, there would several important astrophysical consequences. For instance, there would be a position-dependence in the mean mass to light ratio. This is due to the strong dependence of the mass to light ratio upon the IMF. In R136, this makes the mass density function $\\rho _{{\\scriptsize\\rm m}}$ much different from $\\rho _{{\\scriptsize\\rm l}}~($Malumuth \\& Heap 1994, Brandl et al. 1996) and complicates estimates of the total mass. Padoan et al.'s results indicate that similar effects might occur in spiral galaxies. If the luminosity of a star is taken as $L\\simeq L_{\\odot }m^{y}$, where $y\\simeq 3.5,$ the Jeans expression above would suggest the crude relation $\\propto \\rho ^{3/8}_{{\\scriptsize\\rm l}}$ and yield $\\rho _{{\\scriptsize\\rm m}}\\propto \\rho ^{1+3(1- y)/8}_{{\\scriptsize\\rm l}}\\simeq \\rho ^{0.06}_{{\\scriptsize\\rm l}}$. Unfortunately, previous works have assumed that IMFs are independent of time and position with, specifically, $\\rho _{{\\scriptsize\\rm m}}\\propto \\rho ^{1.0}_{{\\scriptsize\\rm l}}$ throughout a given spiral galaxy (e.g., van Albada et al. 1985). In this {\\it Letter}, surface mass densities of spiral galaxies are computed, for the first time, by explicitly accounting for the possible types of IMF gradients that might exist if theories like those of Padoan et al. are correct. \\par \\noindent {}\\setcounter{section}{2}\\addtocounter{section}{-1} ", + "conclusions": "\\par A direct scaling of R136's IMF to the Galaxy does not dramatically alter the circular velocity curve. However, Models B and C, with their higher yet modest IMF gradients, have nearly flat $v_{\\hbox{{\\scriptsize\\rm circ}}}\\lesssim 220$ km s$^{-1}$ circular velocity curves only before the traditional dark halo component is included. It is interesting to note that if one assumes that these types of models and their $\\sim 10^{1}$-fold mass enhancements are representative of most galaxies, that the fiducial stellar contribution towards the closure density is $\\Omega _{*}\\simeq 0.004~($e.g., Peebles 1993) before accounting for IMF gradients, that the cosmological constant is zero, and that there is no hot dark matter, one would obtain $\\Omega \\simeq \\Omega _{\\hbox{{\\scriptsize\\rm baryon}}}\\simeq 0.04+\\Omega _{\\hbox{{\\scriptsize\\rm gas}}}$, where the closure fraction due to all gas including hot plasma in galactic clusters is $0.007\\lesssim \\Omega _{\\hbox{{\\scriptsize\\rm gas}}}$$\\lesssim 0.08~($Mulchaey et al. 1996). \\par \\medskip Current models of galactic evolution (e.g., Worthey 1994, de Jong 1996) do not account for IMFs that might vary with time and position via the temperature. This is despite prior warnings that the IMF probably has important dependences upon time and position (e.g., Mihalas \\& Binney 1978). In light of the above results, accounting for IMFs with such dependences may be necessary even to obtain results that are only accurate to first order. Accounting for these dependences may, for relatively obvious reasons, clarify our understanding of several astrophysical phenomena including the G-dwarf problem, intrinsic (as a function of radius) and extrinsic (as a function of galactic morphology) metallicity and color gradients, and the Tully-Fisher relation. \\par \\noindent" + }, + "9802/astro-ph9802023_arXiv.txt": { + "abstract": "We report on new {\\it Rossi X-ray Timing Explorer} (RXTE) X-ray spectral analysis of bursts from SGR1806-20, the most prolific SGR source known. Previous studies of bursts from this source noted the remarkable lack of spectral variability both in single bursts as well as from burst to burst. Although we find that the spectrum both within and among bursts is quite uniform we do find evidence for significant spectral changes within bursts as well as from burst to burst. We find that optically thin thermal bremsstrahlung spectra (OTTB) including photoelectric absorption provide the best fits to most bursts, however, other models (power law, Band GRB model) can also produce statistically acceptable fits. We confirm the existence of a rolloff in the photon number spectrum below 5 keV. When modelled as photoelectric absorption and OTTB the inferred column is between $0.8 - 1.2 \\times 10^{23}$ cm$^{-2}$. This value is larger than the $\\approx 0.6 \\times 10^{23}$ cm$^{-2}$ inferred from ASCA observations of the persistent X-ray counterpart, but less than the $\\approx 10.0 \\times 10^{23}$ cm$^{-2}$ indicated by ICE data. ", + "introduction": " ", + "conclusions": "" + }, + "9802/astro-ph9802059_arXiv.txt": { + "abstract": "The X-ray luminosity-temperature relation for nearby $T\\simeq 3.5-10$ keV clusters is rederived using new \\asca\\ temperatures (Markevitch et al.\\ 1998a) and \\rosat\\ luminosities. Both quantities are derived by directly excluding the cooling flow regions. This correction results in a greatly reduced scatter in the $L_X-T$ relation; cooling flow clusters are similar to others outside the small cooling flow regions. For a fit of the form $L_{\\rm bol}\\propto T^\\alpha$, we obtain $\\alpha=2.64\\pm0.27$ (90\\%) and a residual rms scatter in $\\log L_{\\rm bol}$ of 0.10. The derived relation can be directly compared to theoretical predictions that do not include radiative cooling. It also provides an accurate reference point for future evolution searches and comparison to cooler clusters. The new temperatures and $L_X-T$ relation together with a newly selected cluster sample are used to update the temperature function at $z\\sim 0.05$. The resulting function is generally higher and flatter than, although within the errors of, the previous estimates by Edge et al.\\ (1990) and Henry \\& Arnaud (1991, as rederived by Eke et al.\\ 1996). For a qualitative estimate of constraints that the new data place on the density fluctuation spectrum, we apply the Press-Schechter formalism for $\\Omega_0=1$ and 0.3. For $\\Omega_0=1$, assuming cluster isothermality, the temperature function implies $\\sigma_8=0.55\\pm0.03$, while taking into account the observed cluster temperature profiles, $\\sigma_8=0.51\\pm0.03$, consistent with the previously derived range. The dependence of $\\sigma_8$ on $\\Omega_0$ differs from earlier findings, because of our treatment of the slope of the fluctuation spectrum $n$ as a free parameter. For the considered values of $\\Omega_0$, $n=-(2.0-2.3)\\pm0.3$, somewhat steeper than derived from the earlier temperature function data, in agreement with the local slope of the galaxy fluctuation spectrum from the APM survey (Baugh \\& Gazta\\~naga 1996), and significantly steeper than the standard CDM prediction. ", + "introduction": "Clusters exhibit a correlation between their X-ray luminosity and temperature, approximately $L_{\\rm bol}\\propto T^3$ (Mushotzky 1984; Edge \\& Stewart 1991; David et al.\\ 1993; Fukazawa 1997). However, models of cluster evolution that assume no segregation of dark and gaseous matter predict $L_{\\rm bol}\\propto T^2$ (e.g., Kaiser 1986; Navarro, Frenk, \\& White 1995). This discrepancy suggests that the gas is distributed and evolves differently from the dark matter. One of the proposed reasons for this difference is preheating of the intracluster gas by an agent other than gravity, such as supernovae, at some high redshift (Kaiser 1991; Evrard \\& Henry 1991). Significant energy injection to the gas should have been made at the time of its enrichment with heavy elements (e.g., David, Forman, \\& Jones 1991; Loewenstein \\& Mushotzky 1996), although neither the amount nor the epoch of such injection is certain. Models that allow preheating succeed in approximately reproducing the observed $L_X-T$ slope (e.g., Navarro et al.\\ 1995). If the gas was preheated and later subjected to merger shocks, cooler clusters and groups are expected to have a steeper $L_X-T$ dependence than hot clusters (e.g., Cavaliere et al.\\ 1997), for which there is some observational indication (Ponman et al.\\ 1996). Also, preheating should make the redshift evolution of the $L_X-T$ relation slower or absent (Evrard \\& Henry 1991). Indeed, observations exclude any dramatic evolution out to $z\\simeq 0.3-0.5$ (Tsuru et al.\\ 1996; Mushotzky \\& Scharf 1997), although this may also be expected in low $\\Omega_0$ models without preheating (e.g., White \\& Rees 1978; Eke, Navarro, \\& Frenk 1998). The previously reported $L_X-T$ relation at both low and high redshifts exhibits a large scatter which precludes accurate determination of its exact shape and evolution. Fabian et al.\\ (1994) noted that the scatter is mostly due to the strong cooling flow clusters. Cooling flow regions, present in more than a half of all low-redshift clusters (e.g., Edge, Stewart, \\& Fabian 1992), are governed by a different physics than the bulk of the cluster gas, while in extreme cases emitting most of the X-ray luminosity and strongly biasing the wide-beam temperature measurements (e.g., Allen 1998). Cluster evolution models mentioned above predict global properties of the cluster gas and generally ignore the presence of cooling flows, thus it is not obvious that their predictions can be meaningfully compared to the observations. In principle, radiative cooling can be included in the models, but it appears that with greater certainty its effects can be separated in the data, which we attempt in this paper. The $L_X-T$ relation is a necessary tool for derivation of the cluster temperature function using a flux-limited sample. For low-redshift clusters, this function was previously obtained by Edge et al.\\ (1990, hereafter E90) and Henry \\& Arnaud (1991, hereafter HA; see correction in Eke, Cole, \\& Frenk 1996). Since the gas temperature is linked to the total cluster mass, the temperature function provides information on the spectrum of the cosmological density fluctuations (e.g., HA), as well as on $\\Omega_0$ with additional data from higher $z$ (e.g., Henry 1997 and references therein). The least certain link in this line of argument is the conversion from the observed temperature to the cluster mass, which has usually been made assuming cluster isothermality or near-isothermality. However, recent \\asca\\ results (e.g., Markevitch et al.\\ 1998a, hereafter M98) indicate considerable spatial temperature variations which affect the temperature-mass relation. In addition, the temperature function itself is affected by cooling flows via sample selection, individual temperature errors, and the $L_X-T$ relation. All these effects should be corrected for an accurate comparison with theoretical models. In this paper, we use the new \\asca\\ cluster mean temperatures obtained by directly excluding cooling flow regions (M98) and similarly corrected luminosities for a cluster sample selected from the new \\rosat\\ All-Sky dataset. With these data, we rederive the low-redshift $L_X-T$ relation and the temperature function which can be directly compared to theoretical predictions as well as to the oncoming high quality data from higher redshifts. An accurate comparison with models is best performed by detailed simulations and is left for future work; for a qualitative estimate of the constraints that the new data place on the density fluctuation spectrum, we apply a Press-Schechter formalism for $\\Omega_0=1$ and 0.3. We use $H_0=100\\,h$\\kmsmpc; confidence intervals are one-parameter 90\\% unless stated otherwise. ", + "conclusions": "Using the new \\asca\\ cluster temperatures and \\rosat\\ luminosities, both calculated directly excluding the cooling flow regions, we obtain the $L_X-T$ relation in the 3.5--10 keV temperature interval with greatly reduced scatter. This result provides an accurate reference point ({\\em a}) for a comparison with nearby, cooler clusters and groups that should have a different slope if there was preheating, and ({\\em b}) for evolution studies using the oncoming \\axaf\\ data which will provide the necessary angular resolution for cooling flow excision at higher redshifts. The derived $L_X-T$ relation can be directly compared to theoretical and numerical predictions, most of which do not include a detailed treatment of radiative cooling. These new data are also used to rederive the nearby cluster temperature function. This function is interpreted qualitatively by applying the Press-Schechter formalism for $\\Omega_0=1$ and $\\Omega_0=0.3$, to obtain values of $\\sigma_8$ assuming the isothermal and observed cluster temperature profiles (see Table~3). These values should bracket the correct value if some of the observed radial temperature decline in clusters is due to the incomplete thermalization of the gas in the outer regions. The derived slope of the mass fluctuation spectrum at the cluster scale is $n=-(2.0-2.3)\\pm0.3$ for considered values of $\\Omega_0$, which is consistent with the APM galaxy survey at the same scale. Again, our temperature function provides a reference for future accurate evolution studies and for direct comparison with theoretical predictions. It would be interesting to extend the temperature range of this work to lower temperatures, 1--3 keV, where one expects to see, for example, the effects of preheating on the $L_X-T$ relation. It would also greatly strengthen the constraints from the temperature function. Such an extension will be possible when X-ray selected cluster catalogs become available. Since preheating can modify the simple mass-temperature scaling, accurate measurements of cluster masses over a range of temperatures, including the lowest, will be required for the interpretation of such an extended temperature function." + }, + "9802/astro-ph9802090_arXiv.txt": { + "abstract": "We present first results on the Ca isotopic abundances derived from the high resolution Mass Time-of-Flight (MTOF) spectrometer of the charge, element, and isotope analysis system (CELIAS) experiment on board the Solar and Heliospheric Observatory (SOHO). We obtain isotopic ratios $^{40}$Ca/$^{42}$Ca $=$ ($128\\pm47$) and $^{40}$Ca/$^{44}$Ca $=$ ($50\\pm8$), consistent with terrestrial values. This is the first in situ determination of the solar wind calcium isotopic composition and is important for studies of stellar modeling and solar system formation since the present-day solar Ca isotopic abundances are unchanged from their original isotopic composition in the solar nebula. ", + "introduction": "The main motivation to study solar wind composition is to obtain information on the isotopic composition of elements in the Sun. This is important because the Sun constitutes $99.9\\%$ of solar system matter, and for most elements the solar composition could provide the most reliable information on the composition of the primordial solar nebula. Because nowhere in the Sun have temperatures ever been high enough to alter isotopic abundances of heavy elements by nuclear reactions, solar Ca is thought to reflect the original isotopic composition in the solar nebula. Yet, most of what is known of solar isotopic abundances is inferred from terrestrial and from meteoritic abundances with a few exceptions: recently, the isotopic composition of solar wind magnesium, the first analysis of the isotopic abundances of a refractive element in solar matter measured in situ with the spacecraft borne mass spectrometer WIND/MASS, has been reported (\\cite{WindMg}). The solar isotopic composition of the volatile noble gases helium and neon have been determined in situ with the Apollo foil experiments (\\cite{Geiss}). In situ measurements with CELIAS/MTOF (\\cite{Kallenbach}) on board SOHO have confirmed that the solar neon isotopic composition differs significantly from the terrestrial and meteoritic abundances. Isotopic abundances also have been determined for several elements (C, N, O, He, Ne, Mg) in the higher energy (above 10 MeV/amu) solar energetic particles (\\cite{Leske,Mason,Mewaldt,Selesnick}). From similar measurements it has been possible to obtain reliable information on coronal elemental abundances including the abundance of calcium (\\cite{Breneman}). However, the isotopic composition of solar energetic particles (SEP) may not be representative of the solar composition because fractionation processes occur during the particle acceleration and transport and these processes may vary from event to event. The far more fluent solar wind is the most authentic sample of the solar source composition. From recent theoretical models (e.g. \\cite{Bodmer}) it is expected that the isotopic fractionation in the solar wind flow due to differences in Coulomb drag depletes the heavier isotopes by at most a few percent. Therefore measurements of the solar wind isotopic abundances of the heavy element calcium together with previous measurements of magnesium abundances provide additional evidence that the solar isotopic composition of refractive elements agrees with terrestrial and meteoritic values. ", + "conclusions": "From astrophysical and geochemical considerations it can be concluded that the isotopic composition of photospheric Ca must be within small fractions of per mills identical to the terrestrial composition. The similarity of the solar wind results with the terrestrial values suggests the preliminary conclusion that isotopic fractionation within the solar wind plays a minor role. This is supported by the model calculations of~\\cite{Bodmer} that consider differences in the Coulomb drag for isotopes of various elements. Fractionation effects not larger than a few percent are expected for Ne and Mg. For the heavy element Ca even weaker isotopic fractionation has to be expected. Although there are as yet no detailed evaluations of a more extensive time series with calcium isotopic ratios to be determined in different solar wind regimes, we suggest that the Ca abundance discussed in this work is within a few percent of the true solar composition. Table~\\ref{comparison} shows a comparison between the terrestrial, solar and presolar grain calcium isotopic composition. There is no evidence for large variations in these isotopic compositions except for the case of the presolar X-grains that originate from supernovae explosions (\\cite{Hoppe,Nittler}). The large enrichment in $^{44}$Ca in the X-grains is a consequence of the radioactive decay of $^{44}$Ti in supernovae material. \\placetable{comparison}" + }, + "9802/astro-ph9802336_arXiv.txt": { + "abstract": "We present an X-ray spectroscopic study of the prototype far-infrared galaxy NGC6240 from ASCA. The soft X-ray spectrum (below 2 keV) shows clear signatures of thermal emission well described with a multi-temperature optically-thin plasma, which probably originates in a powerful starburst. Strong hard X-ray emission is also detected with ASCA and its spectrum above 3 keV is extremely flat with a prominent iron K line complex, very similar to that seen in the Seyfert 2 galaxy NGC1068 but about an order of magnitude more luminous ($L_{\\rm 3-10keV}\\approx 1.4\\times 10^{42}$\\ergps). The hard X-ray spectrum indicates that only reflected X-rays of an active galactic nucleus (AGN) buried in a heavy obscuration (\\nH$>2\\times 10^{24}$\\psqcm) are visible. This is evidence for an AGN in NGC6240 emitting possibly at a quasar luminosity ($\\sim 10^{45}$\\ergps) and suggests its significant contribution to the far-infrared luminosity. ", + "introduction": "NGC6240 is one of the prototype far-infrared galaxies (FIRGs) that emit most of their bolometric luminosities ($> 10^{11}$\\Ls) in the far-IR waveband (Soifer et al 1984). The origin of the large IR luminosities of this class of objects, comparable to QSOs, has been a big issue since their discovery with IRAS. Either a starburst or an active nucleus embedded in dust is the likely source of the energy output, but which is the major component is still uncertain. X-ray spectoroscopy, particularly with the sensitivity of ASCA in the hard X-ray band, where the obscuration becomes optically thin, can be a powerful tool to probe the energy source. The X-ray properties of several objects belonging to the powerful FIRG class have been investigated with ASCA and the results will appear elsewhere. We present here the results on the brightest object, NGC6240, for which a detailed spectral study is possible. The characteristic properties of powerful FIRGs are found in NGC6240. The 8--1000$\\mu$m luminosity is $L_{\\rm IR}\\sim 6.6\\times 10^{11}$\\Ls (Sanders et al 1989). The distance of the galaxy is assumed to be 100 Mpc throughout this paper, calculated from the cosmological redshift of 0.0245 using the Hubble constant of \\H0 = 75 \\kmpspMpc. Two nuclei separated by 1.8\\sec\\ were found in this galaxy by optical imaging (Fried \\& Schulz 1983) and the distorted galaxy main body is elongated along approximately the N--S direction (PA$\\sim 25^{\\circ}$) with a prominent dust lane, indicating that the system is undergoing a galaxy merger. The large mass of molecular gas, measured with CO line emission, of $M_{\\rm H_2}\\sim 2\\times 10^{10}$\\Ms (Sanders et al 1991; Solomon et al 1997) is confined within the projected size of $\\sim 0.8$ kpc. A deep CO absorption band at $2.3\\mu$m and a powerful K-band continuum ($M_K = -24$) are consistent with the presence of a large number of red giants and supergiants. The reasonably successful modelling of the near-IR emission in terms of starburst stellar population suggests that a starburst is a major power source (e.g., Ridgway, Wynn-Williams \\& Becklin 1994; Shier, Rieke \\& Rieke 1996). Complex filamentary structures of a luminous H$\\alpha$+[NII] nebula ($1.7\\times 10^{42}$\\ergps; extinction-uncorrected) with a total size of 50 $\\times$ 60 kpc, morphologically resembling the well known M82 filaments (McCarthy, Heckman \\& van Breugel 1987) but being much larger in size and $\\sim 30$ times more luminous, were imaged by Heckman et al (1987). The kinematics within the emission-line nebula were studied in detail by Heckman et al (1990) and evidence for a high galactic wind velocity ($\\sim 1000$ \\kmps) has been found. In fact, the optical emission-lines show LINER type excitation (Fosbury \\& Wall 1979; Keel 1990), which is in good agreement with shock heating by the galactic outflow. A very strong H$_2$ 1--0 S(1) emission line at 2.12\\micron ~(e.g., Lester, Harvey \\& Carr 1988; Ridgway et al 1994) that is spatially extended and peaked between the two nuclei (van der Werf et al 1993) also appears to be shock-heated. Although many observations point to a starburst as a major energy source in this galaxy, the ionizing photons deficit problem for Br$\\gamma$ has been pointed out (DePoy et al 1986) and the large [FeII]$\\lambda 1.64$\\micron /Br$\\gamma$ ratio can be reconciled only when an Initial Mass Function (IMF) with a cut-off at $\\sim 25$\\Ms is assumed (van der Werf et al 1993), if one tries to explain the whole IR emission (and ionizing photons) with a starburst alone. Although the nuclear radio source of NGC6240 can be interpreted as powered by a starburst, the ratio of 1.4 GHz radio power and far-IR luminosity is slightly larger than the mean value for spiral and starburst galaxies (Colbert, Wilson and Bland-Hawthorn 1994). Note that the stellar velocity dispersion measured from the CO band, has the highest value so far measured ($\\sim 350$ \\kmps, Doyon et al 1994; Shier et al 1996). The inferred kinematic nuclear mass is too large to match the luminosity (Doyon et al 1994). Bland-Hawthorn, Wilson and Tully (1991) suggested the presence of a massive dark core through a study of kinematics of H$\\alpha$ gas, although the core does not lie at the radio core but at the secondary dynamical system which is located at $\\sim 10$\\sec\\ East of the centre of the galaxy. The ASCA result on NGC6240 has been reported briefly by Kii et al (1997) and Turner et al (1997). We present a detailed spectral analysis using ASCA and ROSAT data and discuss the properties of X-ray emission from a starburst and an active nucleus in this galaxy. ", + "conclusions": "The main results of the ASCA and ROSAT observations of the luminous infrared galaxy NGC6240 can be summarized as follows: \\begin{description} \\item[] (1) The ASCA spectrum at soft X-ray energies ($\\sim$ 0.5--2.5 keV) is rather complex and can be fitted with two thermal components, likely to arise in a strong starburst, which is consistent with the detection of extended X-ray emission in the ROSAT PSPC image ($\\sim$ 0.5--2.0 keV). \\item[] (2) The 3--10 keV spectrum is extremely hard and a strong iron K$\\alpha$ complex feature is present. These results, very similar to those found for the Seyfert 2 galaxy NGC1068, indicate that NGC6240 contains a luminous active nucleus completely obscured by optically thick matter and only visible in reflected light. There has been little direct evidence for a powerful AGN in this object from previous observations at lower energies. \\item[] (3) The inferred luminosity for the AGN component in the hard X-ray band suggests that a significant fraction of the far-infrared luminosity can be powered by the AGN. However, it is not clear whether the same scenario is applicable to ultraluminous infrared galaxies in general. \\end{description}" + }, + "9802/astro-ph9802100_arXiv.txt": { + "abstract": "A detailed analysis of the Fundamental Plane properties of Globular Cluster is performed. If a sample of ordinary King-model clusters is considered, it is found that, in the space ({\\em S-space}) defined by the parameters ({\\rm log} $r_c$, {\\rm log} $\\sigma_0$, $\\mu_V(0)$), their configuration is similar to a straight line. It is shown that, with rather simple assumptions, {\\em a simultaneous explanation of all the observed correlations between S-space parameters} can be provided. It is suggested that, at earlier times, Globular Clusters populated a line in the three-dimensional $S-space$, i.e their original dynamical structure was fully determined by a single physical parameter. \\\\ PACS: 98.20.G ", + "introduction": "The study of self-gravitating stellar systems through the analysis of their distribution into the N-dimensional space of their characteristic parameters has shown to be a very fruitful tool to learn about the formation and dynamical evolution of such objects. The method has been succesfully applied to elliptical galaxies (with the discovery of the so called Fundamental Plane of ellipticals; \\cite{d7}, \\cite{dd87}), to spiral galaxies (\\cite{v84}), to galaxy clusters (\\cite{sh93}), to globular clusters (Djorgovski 1995; hereafter \\cite{D95}) and to all these systems at once (\\cite{bur97}). In this scenario, the globular clusters (GCs) deserve a particular place because they are very simple systems and we also know by direct observations that, at odds with galaxies, they host a single-age/single-metallicity stellar population and they have not been subjected to chemical self-enrichment. Furthermore, correlations between GCs observables can shed some light on the processes that led to formation of globulars, a very important issue by itself (see Meylan \\& Heggie 1997, hereafter \\cite{mh97}, for a complete review). \\cite{D95} showed that Galactic globulars are displaced on a plane into the three-dimensional space defined by the {\\rm log}arithm of their core radii (\\rc, in pc), V-band central surface brightness (\\muv, in mag/arcsec$^2$) and {\\rm log}arithm of the velocity dispersion (\\si in km/s) - hereafter {\\em S-space} - i.e. they constitute a bidimensional manifold in {\\em S-space}. The corresponding scaling law indicates that GCs cores are virialized systems with constant mass to light ratio ($M/L$). In these hypotheses, the generic condition sufficient to obtain the observed Fundamental Plane of Globular Cluster (GCFP) is that the scatter induced in the GCFP by non-homo{\\rm log}y between globulars is smaller than that produced by the dispersion in M/L around a mean value, and by observational errors. This is certainly the most interesting among the three conditions (i.e., virial equilibrium, constancy of M/L and homo{\\rm log}y; see \\cite{ciotti}) but, for the aim of the present analysis, it can be considered as an observational fact. \\cite{D95} showed also that if a space defined by all the photometric, structural and dynamical parameters is considered - hereafter SE-space -, the dimensionality of the manifold is still not more than 3: that means that many of the involved parameters correlates between each other (see also Djorgovsky \\& Meylan 1994 (\\cite{DM94}), \\cite{mh97}). Many of these correlations are not trivial and the origin of most of them is still unclear \\cite{DM94}. However Bellazzini et al. (1996; \\cite{BVFF}) and Vesperini (1997; \\cite{v97}) demonstrated that, for instance, the correlation between integrated cluster magnitude ($M_V$) and the {\\rm log}arithm of the central luminosity density (${\\rm log} ~\\rho_0$) is very likely to have been settled at the time of GC formation, and that dynamical evolution, either intrinsic or due to the interaction with the host galaxy, is much more efficient in {\\em damaging} any existing correlation between structural parameters than in settling one such correlation by means of evolutionary or selective disruption effects. These latter studies showed that it is possible to recover informations on the initial dynamical/structural conditions of a GCs system from present-day correlations between {\\em SE-space} parameters. In this short note I present a number of hints suggesting that Galactic GCs were formerly displaced on a Straight Line in the {\\em S-space}, i.e. they formed a one-dimensional manifold in such space. Furthermore I present a simple and general explanation for all the observed mono-variate correlations between \\s parameters, comprehensive of a new, more satisfactory, interpretation of the correlation between cental surface brightness and velocity dispersion. ", + "conclusions": "The distribution of the Galactic KMCs into the \\s is very similar to a straight line lying onto the Fundamental Plane of Globular Clusters. I have shown that {\\em the simple assumptions }(supported by observations) {\\em $M/L \\simeq const.$ and $M_c \\sim const.$ provide a simultaneous explanation for all the three mono-variate correlations present in the data} (and also for GCFP correlations). In particular this provides a new and more satisfactory interpretation of the correlation between \\muv and \\lsi. Some interesting by-products have also been obtained: {\\em a)} the thickness of the GCFP and of FPEG at any mass are very similar; {\\em b)} the range of M/L covered by GCs is constrained to be significantly less than an order of magnitude with a fully {\\em model-independent} procedure; {\\em c)} the current definition of observables, coupled with a strong $M/L \\simeq const.$ constraint, create a condition in which even if the $M_c$ range of GC would encompass 4 order of magnitude, we would still observe some significant correlation in the planes (\\lrc -- \\muv) and (\\muv -- \\lsi). In my opinion, the most far reaching conclusion of the present analysis is that globulars {\\em were distributed along a Fundametal Straight Line in S-space} at early times. This conclusion provides a very useful constraint for models of globular cluster formation, i.e. globulars were born with nearly {\\em the same core mass} (or rapidly settled to this status), within a range that was probably much narrower than the one presently observed. Present knowledge of the formation of globular clusters is rather poor (see \\cite{fr85}, \\cite{vp95}, \\cite{hp94}, \\cite{mh97}) and any observational constraint on initial conditions of this system has to be considered very valuable. George Djorgovski is warmly thanked for having introduced me to the mysteries of Principal Component Analysis. I am indebted to Paolo Montegriffo, Enrico Vesperini, Stefano Sandrelli and Flavio Fusi Pecci for many useful discussions. Very special thanks are owed to Luca Ciotti for many insightful discussions and a critical reading of draft versions of the manuscript. Mrs Paola Ballanti is warmly thanked for her professional expertise in english language. \\clearpage" + }, + "9802/astro-ph9802270_arXiv.txt": { + "abstract": "We analyze the window functions for the spherical harmonic mode estimators of all--sky, volume limited surveys considering evolutionary effects along the past light--cone which include the deviation of the distance scale from a linear relationship with redshift, linear peculiar velocity corrections, and linear evolution of the density perturbations. The spherical harmonic basis functions are considered because they correspond most closely to the symmetries of typical survey geometries and of the light--cone effects we consider. Our results show substantial broadening of the windows over that expected by ignoring light--cone effects, indicating the difficulty of measuring the power spectrum independently from cosmology. We suggest that because of light--cone effects, deep redshift surveys should either be analyzed in conjunction with CMBR data which determines the cosmological parameters, or by using a Bayesian likelihood scheme in which varying cosmological parameters and a simple parameterization of the primordial power spectrum are assumed as the priors, so that observed data can be mapped from redshift to real space. The derived power spectrum can then be compared to underlying models of fluctuation generation and growth in structure formation to evaluate both these models and the cosmological priors. ", + "introduction": " ", + "conclusions": "" + }, + "9802/astro-ph9802046_arXiv.txt": { + "abstract": "We have analyzed 0.35--7.5 keV X-ray spectra of the center of the Perseus cluster collected using the Broad Band X-Ray Telescope (BBXRT) on the Astro-1 mission. These spectra are particularly useful for examining the nature of the X-ray absorber in cooling flows because of BBXRT's sensitivity between 0.35 and 1.0 keV. We confirm that there is X-ray absorption above that expected from gas in our own galaxy. Further, the absorbing medium is deficient in helium. However, the energy of the K-edge of oxygen is consistent with neutral material (at the redshift of the cluster) and is not consistent with any ionized state of oxygen. It is not possible to completely ionize helium and have oxygen neutral so the apparent helium deficiency cannot be due to ionization. We propose that the X-ray absorption is due to dust grains that have condensed out of a medium in which helium remains ionized. This model satisfies all the observational constraints but is difficult to understand theoretically. ", + "introduction": "Clusters of galaxies are prodigious sources of X-ray emission, produced by a $10^7-10^8$~K plasma confined by the gravitational field of the cluster. In the inner tens of kpc of many clusters the radiative cooling time of this gas is less than the age of the cluster implying that the gas cannot be static and must be cooling and flowing into the center (see eg Fabian 1994 and references therein). This picture has been verified by both X-ray imaging and spectroscopy with the implied cooling rates being, in some cases, many hundreds of Solar masses per year. This led to the central mystery of cooling flows : if so much mass is cooling through the X-ray band ($T\\sim 10^7$K) what happens to it after that ? Solutions can be roughly classified into three camps : a) the X-ray data analysis is wrong, b) there is energy input which prevents the gas cooling out of the X-ray regime, and c) the gas cools down and condenses into something that is not observed (eg low mass stars). Solution a) has been shown to be incorrect by the consistent results obtained using a variety of different detectors and analysis methods. Solution b) is much argued over and seems to us to require an unlikely amount of fine tuning. Solution c) is unsatisfactory in that it posits placing the mass in something that cannot be seen. There is some evidence for star formation in the galaxies at the centers of cooling flows but with a conventional initial mass function this can mop up only 1--10\\% of the available mass (e.g. McNamara 1997). Solution c) received a large boost following the observation by White \\etal (1991) that cooling flows showed X-ray absorption in excess of that from our own Galaxy. The column densities observed were consistent with the accumulation of cooling flow matter over the lifetime of the cluster. These excess absorptions were confirmed by analysis of spectra obtained using ROSAT (Allen \\etal 1993) and ASCA (Fabian \\etal 1994). However, this cold, absorbing material is not seen in HI or CO (e.g. O'Dea \\& Baum 1997). Again, we obtain a result using X-ray data but cannot find the corollaries in other wavebands. A further problem is that the excess absorption does not always show up in the analysis of spectra obtained using the ROSAT PSPC (e.g. Sarazin 1997). Since the ROSAT PSPC has a lower energy cut-off than ASCA it should provide a more sensitive measurement of absorption. In this paper we report investigations on the nature of the X-ray absorber in the Perseus cluster, a much-studied, massive cooling flow centered on the galaxy NGC 1275. We use data from the Broad Band X-Ray Telescope (BBXRT) that flew on the Astro-1 Space Shuttle mission December 2--10, 1990. BBXRT used telescopes similar to those on ASCA but with segmented Si(Li) solid-state detectors instead of CCDs. The BBXRT Si(Li) detectors have slightly lower spectral resolution than the ASCA CCDs but they do not have the instrumental oxygen absorption which reduces the CCDs' efficiency immediately above 0.5 keV. Since the primary absorption feature in the energy range covered by these detectors is the O K edge, BBXRT data provide a more sensitive measurement of the redshift and detailed shape of the X-ray absorption than is possible with ASCA. ", + "conclusions": "We have shown that the BBXRT spectra of the Perseus Cluster are consistent with a cooling flow absorber at the redshift of the cluster. The data are marginally inconsistent with the absorber being local to our galaxy. More important, we have demonstrated that the absorber is deficient in helium (and, with less confidence, neon). The only reasonable explanation of this is that the absorber is comprised of dust grains formed out of gas in which helium is completely ionized. The absorption model we have fit assumes that the dust grains are a screen in front of the cooling flow. This is oversimplistic and it is likely that they are distributed throughout the cooling flow. The present data are not good enough to explore this issue. A model in which cold grains are distributed throughout the cooling region and are heated only where there are emission line filaments satisfies all the observational constraints. However, we do not understand how to keep the oxygen on the grains, how to prevent the grains being destroyed, or how to keep the grains at very low effective temperatures. BBXRT also observed the cooling flow clusters Abell 262 and Abell 496. These data are not as high quality as those for Perseus and all we can conclude for these clusters is that there is an O K edge present whose depth exceeds that expected from the Galactic column (Mackenzie, Schlegel \\& Mushotzky 1996). Our result could be confirmed using the the backside-illuminated ACIS CCDs on AXAF. In the future, high resolution spectroscopy will provide detailed information about abundances, ionization states, and grain composition from the energies, depths, and detailed shapes of absorption edges (see e.g. Woo, Forrey \\& Cho 1997). We thank Ari Laor for reminding us that we had never published this data. We thank Hagai Netzer and Tim Kallman for helpful discussions. This research has made use of data obtained through the High Energy Astrophysics Science Archive Research Center Online Service, provided by the NASA/Goddard Space Flight Center. \\newpage" + }, + "9802/gr-qc9802015_arXiv.txt": { + "abstract": "We examine the field equations of a self-gravitating global string in low energy superstring gravity, allowing for an arbitrary coupling of the global string to the dilaton. Massive and massless dilatons are considered. For the massive dilaton the spacetime is similar to the recently discovered non-singular time-dependent Einstein self-gravitating global string, but the massless dilaton generically gives a singular spacetime, even allowing for time-dependence. We also demonstrate a time-dependent non-singular string/anti-string configuration, in which the string pair causes a compactification of two of the spatial dimensions, albeit on a very large scale. ", + "introduction": "Topological defects crop up widely in physics; in cosmology they have been put forward as a possible source for the density perturbations which seeded galaxy formation \\cite{BSV}. Phase transitions in the early universe can give rise to various forms of topological defect. A defect is a discontinuity in the vacuum and is classified according to the topology of the vacuum manifold of the field theory model being used. Disconnected vacuum manifolds give domain walls, non-simply connected manifolds, strings, and vacuum manifolds with non-contractible spheres give monopoles. Strings and monopoles can be further divided into local and global defects depending on the nature of the symmetry broken. Local defects are formed when the symmetry broken is gauged, and the Higgs mechanism typically removes any Goldstone bosons, meaning that the defect has a well-defined core and finite energy per unit defect area. Global defects on the other hand (with the exception of the domain wall which does not result from the breaking of a continuous symmetry) typically have a residual Goldstone field in the Lagrangian, which translates to a power law fall-off in the core energy density and a divergent total energy. More unusual defects, such as semi-local strings \\cite{VA} also exist, but we do not consider these here. Given this difference between local and global defects, we might expect them to have significantly different behaviour, and this is particularly evident in their coupling to gravity. Whilst local strings \\cite{LSTR} and monopoles \\cite{LMON} are well-behaved and asymptote flat or locally flat spacetimes, global strings \\cite{GSTR,G} and monopoles \\cite{BV,HL} have strong effects at large distances. Indeed, the spacetime of a global string was for some time thought to be singular \\cite{GSTR}. In fact, the spacetime is time-dependent \\cite{G} with a de-Sitter expansion along the spatial extent of the defect. The spacetime is also characterised by an event horizon at a finite distance from the string core, although this horizon appears to be unstable to perturbations \\cite{INSTAB}. Of course, this work has all been performed within the context of general relativity. At sufficiently high energy scales it seems likely that gravity is not given by the Einstein action, but becomes modified, for example, by superstring terms which are scalar-tensor in nature \\cite{LESG}. Low energy string gravity is reminiscent of the scalar-tensor theories of Jordan, Brans and Dicke \\cite{JBD}, and is equivalent to Brans-Dicke theory for a particular value of the Brans-Dicke parameter : $\\omega=-1$. The implications of Brans-Dicke gravity for defects have been explored for local strings \\cite{GO} and global monopoles \\cite{BR}. However, since the Brans-Dicke parameter is known to be constrained by $\\omega > 500$ \\cite{REAS}, we are more interested in exploring the low energy superstring action, particularly with a massive dilaton, although massless dilatons will provide a comparison with Brans-Dicke results. Moreover, some recent superstring inspired models of inflation \\cite{AR} appear to allow for cosmic string defects to form of either the local or global variety, therefore it is useful to know how these objects interact with the dilaton. In this paper we will consider the implications of superstring gravity for global cosmic strings. In particular we will address the question of whether a non-singular spacetime exists for the string. Recently Sen et.\\ al.\\ \\cite{SEN} studied global strings in Brans-Dicke theory. We disagree with their claim that a static, non-singular solution exists for the string spacetime. The structure of this paper is as follows. First, we briefly review the global string in Einstein gravity. We then study the global string in superstring gravity, for both massless and massive dilatons. Finally, we compare our results with the literature and present our conclusions - in particular, we present a novel solution which represents a self-gravitating string/anti-string pair on a closed two-dimensional spatial section. ", + "conclusions": "We have studied the behaviour of the metric and dilaton field of a global cosmic string in superstring gravity for an arbitrary coupling of the string Lagrangian to the dilaton : $e^{2a\\phi} {\\cal L}$. For both massless and massive dilatons, we have demonstrated the existence of a non-singular spacetime for the string if we include an exponential expansion along the length of the string $e^{b_0t}$, with $b_0 > 0$. In addition we have the further restriction that $a=-1 - 2 \\epsilon K_2$ for the massless dilaton. In both cases, the spacetime is characterised by an event horizon at finite distance from the string core. Near this point, the asymptotic solution for the metric is \\be ds^2 \\approx b_0 (r_0-r)^2 \\left [ dt^2 -\\cosh^2 \\sqrt{b_0}t dz^2 \\right ] - dr^2 - C_0^2 d\\theta^2 \\ee However, since this non-singular solution is very similar to the Einstein global string, we expect that it too will be unstable. The metric at intermediate points will be given by the Cohen-Kaplan \\cite{GSTR} solution, and the dilaton by either (\\ref{massless}) if it is massless, (\\ref{massive}) if it is massive and $a\\neq-1$, or (\\ref{massivea}) if the dilaton is massive and $a =-1$. Thus we expect that the cosmological effects of global strings deriving from their purely gravitational or metric properties \\cite{HH} will be little altered from the Einstein case. The main difference will be dilaton production by global strings. In this case, as with the global monopole \\cite{DG}, the Damour-Vilenkin \\cite{DV} bounds on the dilaton mass hold for $a\\neq-1$, but if a global string couples to a massive dilaton with $a = -1$, then the Damour-Vilenkin bound is weakened slightly. For example, for a TeV mass dilaton, the Damour-Vilenkin limit on the symmetry breaking scale $\\eta$ is 10$^{13}$ GeV, which is weakened to 10$^{14}$ GeV in the case $a = -1$. In the course of our analysis, we have shown that the spacetime of a static ($b_0=0$) global string in dilaton gravity is necessarily singular. This is in disagreement with the work of Sen {\\it et.\\ al.\\ } \\cite{SEN} who studied static global strings in Brans-Dicke theory, apparently finding non-singular solutions to the far-field equations. By utilising the dynamical system approach, we can see that while quite valid as solutions to the far-field equations, these are not solutions to the full field equations. Massless dilatonic gravity corresponds to Brans-Dicke theory for the particular parameter values $\\omega=-1$ and $a=0$. With $b_0=a=0$ the far-field dynamical system (\\ref{twods}) is \\bea \\dot{f} & = & f^2 - \\frac{5}{2}g^2 \\\\ \\dot{g} & = & 2 f^2 - \\frac{5}{2} g^2 - fg \\eea with a single fixed point at $f=g=0$ (see Figure \\ref{fig:plot2}). The separatrices are $f-g$ and $f+ \\frac{\\sqrt{5}}{2}g$. The non-singular solution found by Sen {\\it et.\\ al.\\ } corresponds to \\be f=g = \\frac{2}{3(r-r_0)} \\ee where $r_0 \\le 0$. This is the part of the separatrix $f=g$ lying in the upper-right quadrant of the phase plane with the flow towards the origin. However, $f-g = {\\dot A}$, which is zero at the core of the global string, and integrating (\\ref{mfe2}) shows \\be (f-g) Ce^A = - {\\epsilon\\over4} \\int_0^\\rho Ce^{3A} e^{4\\phi} (X^2-1)^2 d\\rho <0 \\ee i.e.\\ $g$ is strictly greater than $f$ outside the core. Thus the trajectory $f=g$ cannot correspond to the exterior of a physical global string. Finally, we would now like to comment on the possibility discussed earlier, namely that $C\\to 0$ at some finite $r_1$, giving a solution on a compact $(r,\\theta)$ section. Since the field equations are symmetric under the transformation $r \\to r_1-r$, their solution must also be symmetric. Thus at $r_0= r_1/2$, we have $C'=A'=\\phi'=0$. Defining $f,g$ and $h$ as before, this means that $f=g=h=0$ at $r_0$. Whether the dilaton is massive or massless, we can reduce the far-field equations to the same two-dimensional dynamical system (\\ref{twods}). The symmetric solution must correspond to the trajectory going through the origin, since $f=g=0$ at $r_0$. For $f,g>0$ this trajectory is trapped between the non-singular isolated string trajectory and the line $y=(1+K_2 \\epsilon)x$. Hence, we can again argue that for some $b_0>0$ the solution represented by this trajectory matches onto the core of the global string near $r=0$. For $f,g<0$, we can similarly argue that as $r\\to r_1$ and $f,g\\to -\\infty$, the trajectory matches onto the core of an anti-string located at $r_1$. The $(r,\\theta)$ spatial sections of this spacetime are compact, and are qualitatively depicted in Figure \\ref{fig:plot3}, with the string and anti-string at opposite poles. We can estimate the scale at which this compactification occurs by integrating (\\ref{mfe1}), which gives $r_0$ of the order of $e^{1/\\epsilon}$. For $\\epsilon \\sim 10^{-6}$, this is way beyond the current cosmological horizon, however, it is tempting to extrapolate this solution to larger values of $\\epsilon$. If $\\epsilon$=O(1), then the $(r,\\theta)$ sections would be compact at a scale of order $\\epsilon$, and the spacetime would be effectively two-dimensional with an exponential expansion in the spatial dimension. Of course this is way beyond the validity of our analysis, however, it is interesting to consider in the light of topological inflation scenarios \\cite{TOP}. In these, a Planck scale topological defect is considered as a source for inflation. If this string/anti-string solution persists at high energy, then the global string would not be a suitable candidate for topological inflation. \\vskip 1cm \\noindent{\\it Note added} After completing this work, we note that Boisseau and Linet \\cite{BL} have recently computed the exterior metric of a global string in Brans-Dicke gravity. \\vskip 2cm" + }, + "9802/astro-ph9802264_arXiv.txt": { + "abstract": "We report a $3\\sigma$ upper limit of $220$\\,$\\mu$Jy on any 22-GHz continuum emission coincident with the central engine in NGC 4258. If NGC 4258 is powered by an advection-dominated accretion flow, this radio upper limit implies that the inner advection-dominated flow cannot extend significantly beyond $\\sim10^{2}$ Schwarzschild radii. ", + "introduction": "NGC 4258 is a weakly active Seyfert 2 galaxy possessing a highly obscured central X-ray source with a 2--10 keV luminosity of $4\\times10^{40}$ \\ergs\\ (Makishima \\etal\\ 1994) and nuclear continuum and narrow line emission seen in reflected, polarized optical light (Wilkes \\etal\\ 1995). The galaxy also harbors one of the first known nuclear megamasers (Claussen, Heiligman, \\& Lo 1984), and VLBA observations reveal a nearly edge-on, extremely thin, slightly warped Keplerian disk (Watson \\& Wallin 1994; Greenhill \\etal\\ 1995; Miyoshi \\etal\\ 1995; Moran \\etal\\ 1995; Herrnstein, Greenhill, \\& Moran 1996). The masers extend from 0.13 to 0.26 pc (for a distance of 6.4 Mpc), and the Keplerian rotation curve requires a central binding mass, $M$, of $3.5\\pm0.1\\times10^{7}$ \\msun\\ within 0.13 pc. The velocity centroid of the disk agrees well with the optically-determined systemic velocity of the galaxy, and the rotation axis of the disk is aligned with the inner portion of large-scale twisted jets seen in radio to X-ray emission (Cecil \\etal\\ 1992). VLBA continuum observations also reveal a subparsec-scale jet oriented along the disk axis (Herrnstein \\etal\\ 1997). The central mass estimate provided by the maser observations corresponds to an Eddington luminosity, $L_{E}$, of $4.4\\pm0.1\\times10^{45}$\\,\\ergs. The bolometric luminosity, $L$, is more difficult to estimate because the central edge-on disk obscures the nuclear emission. Wilkes \\etal\\ (1995) estimate $L \\sim 10^{42-44}$ \\ergs\\ based on observations of the nuclear continuum in optical, polarized light. However, the 2--10 keV X-ray flux suggests $L\\sim4\\times10^{41}$ \\ergs, based on the argument that the X-ray luminosity typically accounts for $\\sim10$\\% of $L$ in AGN (Mushotsky, Done, \\& Pounds 1993). Hence $L\\sim10^{42\\pm1}=10^{-3.6\\pm1}L_{E}$, and the central source in NGC 4258 is highly sub-Eddington. Between 0.13 and 0.26 pc ($4\\times10^{4}$ and $8\\times10^{4}$ Schwarzschild radii, \\rg), the masers appear to trace a cool, thin accretion disk: the temperature in the maser layer must be between approximately 300 and 1000\\,K to support maser action, and the aspect ratio of the layer is less than 0.3\\% (Moran \\etal\\ 1995). Unfortunately, the structure within 0.13 pc is not directly observable, and the precise nature of the NGC 4258 central engine remains obscure. One possibility is that the outer thin disk traced by the masers extends to the central black hole, and that the NGC 4258 central engine is fueled by an optically thick, geometrically thin accretion disk (Shakura \\& Sunyaev 1973). Since the radiative efficiency, $\\eta$ (defined through $L=\\eta\\mdot c^{2}$, where \\mdot\\ is the accretion rate), of such a disk is high, the sub-Eddington luminosity of NGC 4258 then implies a correspondingly sub-Eddington \\mdot. Specifically, for $\\eta=0.1$, $\\mdote\\equiv\\mdot/\\mdotedd=10^{-3.6\\pm1}$. Here \\mdotedd\\ is the Eddington accretion rate, given by $2.2\\times10^{-8}M$\\,\\smy. Neufeld \\& Maloney (1995) argue that for $\\mdote\\sim10^{-4.0}\\alpha$, the outer disk, which is obliquely irradiated by the central X-ray source as a result of the warp, changes from cool molecular gas to warm atomic gas at about 0.23 pc, providing a natural explanation for the observed outer edge to the maser emission. Here, $\\alpha$ is the standard Shakura-Sunyaev parameterization of the kinematic viscosity. In the Neufeld \\& Maloney model, $\\eta\\sim10^{-0.6\\pm1}\\alpha^{-1}$. Alternatively, Lasota \\etal\\ (1996; hereafter L96) have proposed that NGC 4258 harbors an optically thin advection-dominated accretion flow (ADAF) at small radii, within the cool molecular disk. Narayan \\& Yi (1994, 1995a\\&b; hereafter NY95b) have demonstrated the stability and self-consistency of a geometrically thick, optically thin accretion flow in which an extremely hot ($T\\sim10^{12}$\\,K) ion plasma coexists with much cooler ($T\\sim10^{9.5}$ \\,K) electrons. For the ions, the radiative timescale is much longer than the accretion timescale. Thus, in the case of accretion onto a black hole (BH), the majority of the viscously dissipated energy is carried through the event horizon of the BH by the hot ions, and the ADAF radiative efficiency is several orders of magnitude less than that of the standard thin disk. The appeal of the two-temperature ADAF models is that the optically thin synchrotron, bremsstrahlung, and Compton emission from the $\\sim10^{9.5}$\\,K electrons can account for the spectra of a broad variety of accreting systems over many decades in energy (see Narayan 1997 for a recent review). For example, efforts to model the radio-to-$\\gamma$-ray spectrum of the galactic source Sgr A$^{*}$ with an ADAF spectrum are reasonably successful (Narayan, Yi, \\& Mahadevan 1995; Manmoto, Mineshige, \\& Kusunose 1998; Narayan \\etal\\ 1998). This is an important result: optically thick, thin disk models emit largely as blackbodies and are generally too cool to account for the X-ray emission observed in accreting binaries and active galactic nuclei (AGN). The two-temperature ADAF solution is valid for accretion rates less than $\\mdotc\\sim1.3\\alpha^{2}$ (Esin \\etal\\ 1997) or, equivalently, for bolometric luminosities less than $\\sim1.3\\alpha^{2}L_{E}$. NGC 4258 satisfies this requirement for $\\alpha>10^{-1.9\\pm0.5}$. In general, the ADAF models are parameterized in terms of $M$, $\\alpha$, $\\mdote$, $\\beta$ (the ratio of gas pressure to total pressure), and \\rtrans\\ (the transition radius from the outer thin disk to the inner ADAF). The NGC 4258 ADAF model of L96 is constrained primarily by the maser disk binding mass and the 2--10 keV X-ray luminosity and spectral index, and secondarily by the polarized optical flux. L96 assume $\\beta=0.95$ and find that an NGC 4258 ADAF requires $\\mdote\\sim10^{-1.8}\\alpha$. More recent calculations suggest $\\mdote\\sim10^{-1.6}\\alpha$ (Narayan, personal communication), and the condition $\\mdote\\la\\mdotc$ requires $\\alpha\\ga0.02$ in NGC 4258. In this model, $\\eta\\sim10^{-3\\pm1}\\alpha^{-1}$, and the NGC 4258 ADAF is very rapid, but very inefficient in producing radiation. In general, the hot ADAF electrons are expected to generate significant radio emission via the thermal synchrotron process. Unfortunately, most of the underluminous AGN and quiescent ellipticals which plausibly harbor ADAFs also possess core-jet radio structures that are not directly associated with the ADAF itself. In practice it is extremely difficult to discriminate this emission from the ADAF synchrotron emission. Even when VLBI provides sufficient angular resolution to resolve individual components within the radio core, it is usually impossible to determine which, if any, of this emission originates precisely at the central engine, where the ADAF emission must arise. Because the relative position of the center of mass of the sub-parsec molecular disk can be measured to a fraction of a milliarcsecond with VLBI, NGC 4258 provides a rare opportunity to circumvent the ambiguity associated with core-jet emission, and to test the radio portion of the proposed ADAF spectrum directly. ", + "conclusions": "\\label{s:adaf.3} The most direct method to discriminate between the standard thin-disk and ADAF accretion modes would be to determine the accretion rate through the outer thin disk. For a Shakura-Sunyaev accretion disk in hydrostatic equilibrium (c.f. Frank, King, \\& Raine 1992) \\bq \\frac{\\mdote}{\\alpha}\\simeq 530\\frac{\\dmass c_{s}^{2}}{M^{3/2}(r_{2}^{1/2}-r_{1}^{1/2})} ~~, \\label{eq:danger} \\eq where \\dmass\\ is the disk mass (in solar masses) between $r_{1}$ and $r_{2}$ (in pc), and $c_{s}$ is the isothermal sound speed in \\kms. The observed upper limit in the thickness of the maser layer, together with the mere presence of the masers, implies $c_{s}\\la2.5$\\,\\kms\\ (or $T \\la 1000$K) (Moran \\etal\\ 1995). The simplest assumption is that the masers lie in the midplane of the disk and that the midplane density must be less than $\\sim10^{10}$\\,cm$^{-3}$ to avoid thermal quenching of the population inversion. In this case, $\\dmass\\la5.8\\times10^{4}$\\,\\msun\\ between 0.13 and 0.26 pc and $\\mdote/\\alpha\\la10^{-2.2}$. A second, more observationally motivated constraint on \\dmass\\, comes from the precision of the Keplerian rotation curve of the high-velocity masers. While the statistical scatter in the LOS velocities around the best-fitting Keplerian rotation curve is about 3\\,\\kms, the systematic deviation from Keplerian rotation (\\dvsys) across the masing zone may be considerably larger due to uncertainties in the fitted-disk parameters, which are partially correlated. A formal \\chisq\\ analysis demonstrates $\\dvsys\\la20$\\,\\kms. This requires $\\dmass\\la1.5\\times10^{6}$\\,\\msun\\ and $\\mdote/\\alpha\\la10^{-0.8}$. This is about six times larger than the ADAF accretion rate. Thus, the maser data do not definitively rule out either mode of accretion in NGC 4258, and the present continuum observations provide an important additional constraint on the models. We associate the center of the maser disk with the NGC 4258 central engine, and interpret the continuum non-detection in the context of the ADAF models using the analytical expressions presented in Mahadevan (1997; hereafter M97), which duplicate reasonably well the numerical results of NY95b. M97 includes all the relevant electron heating and cooling terms, but considers only the region between 3 and $10^{3}$ \\rg, assumes that the electron temperature, $T_{e}$, is constant and much less than the ion temperature over these radii, and treats the flow as purely spherical. These approximations are justified by the more detailed analysis of NY95b. M97 treats the ADAF as a relativistic thermal plasma, which is valid for $\\mdote\\ga10^{-4}\\alpha^{2}$ (Mahadevan \\& Quataert 1998) and is justified in NGC 4258. M97 shows that the radius, $r_{\\nu}$, at which the spherical flow becomes optically thick to synchrotron self-absorption depends on frequency, $\\nu$, as \\bq r_{\\nu}=0.47T_{e}^{8/5}\\nu^{-4/5}M^{-2/5}\\mdote^{3/5}\\alpha^{-2/5}\\mbox{\\, \\rg}, \\label{eq:maha1} \\eq where we have assumed $\\beta=0.5$. To a very good approximation, the synchrotron spectrum, $L_{\\nu}$, is completely determined by the emission at $r_{\\nu}$, which may be approximated by a blackbody spectrum, $B_{\\nu}$, in the Rayleigh-Jeans limit (M97): \\bq L_{\\nu}=\\pi B_{\\nu}4\\pi r_{\\nu}^{2}=2.3\\times10^{-25}T_{e}^{21/5}\\nu^{2/5}M^{6/5}\\mdote^{6/5}\\alpha^{-4/5}\\mbox{ \\ergs}. \\label{eq:maha2} \\eq Equations~\\ref{eq:maha1} and \\ref{eq:maha2} show that both $r_{\\nu}$ and $L_{\\nu}$ are steep functions of the electron temperature. However, the analytical treatment of M97 and the more rigorous numerical work of NY95b both indicate that for high accretion-rate systems, the electron temperature is confined to a narrow range, and is essentially independent of the central mass. Specifically, M97 finds that for $\\mdote\\sim10^{-1.5}\\alpha$, $2.1\\times10^{9} \\la T_{e} \\la 3.1\\times10^{9}$ over 9 decades in $M$. For $\\mdote=10^{-1.6}\\alpha$, $M=3.5\\times10^{7}$\\,\\msun, and $T_{e}=2.15\\times10^{9}$\\,K, equation~\\ref{eq:maha2} indicates that $F_{22}\\ga220$\\,$\\mu$Jy for $\\alpha>0.009$. Here $F_{22}$ is the 22-GHz ADAF flux density assuming a distance of 6.4\\,Mpc. For $\\alpha\\la10^{-2}$, $\\mdote\\ga\\mdotc$ and advection is not a viable solution. Hence, the ADAF models significantly overestimate the actual 22-GHz emission over the full range of allowed $\\alpha$. Here, we have assumed that \\rtrans\\ is larger than the effective 22-GHz photosphere (\\rtt) as given by equation~\\ref{eq:maha1}, a necessary prerequisite to the assumption of constant $T_{e}$. However, the absence of detectable 22-GHz continuum emission in NGC 4258 may be consistent with an ADAF that is truncated such that $\\rtrans<\\rtt$. Equation~\\ref{eq:maha1} indicates that $\\rtt\\simeq300\\alpha^{1/5}$\\,\\rg\\ for the NGC 4258 ADAF, suggesting that $\\rtrans\\la10^{2}$\\,\\rg\\ for typical values of $\\alpha$. A second, more stringent upper limit on \\rtrans\\ comes from assuming that, for those frequencies below which the condition $r_{\\nu}\\ga\\rtrans$ is satisfied, the emission follows the $T=2.5\\times10^{9}$ blackbody spectrum of the outermost ADAF shell at $r=\\rtrans$. In this case, the 22-GHz upper limit requires $\\rtrans\\la80$\\,\\rg. {\\it Thus, the 220\\,$\\mu$Jy, 3$\\sigma$ upper limit on any compact 22-GHz emission coincident with the central engine in NGC 4258 implies that the proposed central ADAF cannot extend significantly beyond $\\sim10^{2}$\\,\\rg}. More rigorous numerical methods also lead to ADAF solutions consistent with the present 22-GHz upper limit for $\\rtrans\\la10^{2}$\\,\\rg\\ (Narayan, personal communication). In addition to affecting the radio emission, the transition radius affects both the hard X-ray spectrum and the optical/UV emission. In the former case, the ADAF is Compton cooled by the soft photons from the outer cool disk, and as \\rtrans\\ moves inward, the hard X-rays are suppressed. L96 find that $\\rtrans\\ga10$\\,\\rg\\ is consistent with the Makishima \\etal\\ (1994) measurements. The optical/UV spectrum depends on \\rtrans\\ because the blackbody spectrum of the non-molecular portion of the disk peaks at these wavelengths. Wilkes \\etal\\ (1995) have measured the 5500 \\AA\\ flux toward NGC 4258 in polarized light. This emission presumably arises in the central engine and has been scattered into our LOS. The corresponding 5500 \\AA\\ central engine luminosity is highly sensitive to the type of scattering screen invoked (dust or electrons), and Wilkes \\etal\\ (1995) estimate that $L_{5500}\\sim10^{37-39}$\\,\\ergs\\,\\AA$^{-1}$. L96 find that the NGC 4258 ADAF contributes only about $10^{36.3}$\\,\\ergs\\,$\\AA^{-1}$\\ at 5500\\,\\AA, and here we estimate the contribution of the thin disk to the 5500\\,\\AA\\ luminosity for a range of \\rtrans. The inferred isotropic luminosity of an optically thick, steady, thin disk truncated at \\rtrans\\ is (c.f. Frank, King, \\& Raine 1992): \\bq L_{\\lambda}^{(td)}\\simeq\\frac{16\\pi^{2}h\\nu^{2}\\cos{i}}{\\lambda^{3}}\\int_{\\rtrans}^{\\mbox{$r_{out}$}}\\frac{rdr}{e^{h\\nu/kT(r)}-1}\\mbox{ \\ergs\\,\\AA$^{-1}$}, \\label{eq:fkr1} \\eq where $T(r)=2.2\\times10^{5}\\mdot_{26}^{1/4}M_{8}^{1/4}r_{14}^{-3/4}$ K, $i$ is the angle between the disk normal and the LOS ($\\sim82^{\\circ}$ in NGC 4258), $\\mdot_{26}$ is the accretion rate in units of $10^{26}$\\,g\\,s$^{-1}$, $M_{8}$ is the central mass in units of $10^{8}$\\,\\msun, and $r_{14}$ is the radius in units of $10^{14}$\\,cm. For $\\mdot/\\alpha=10^{-1.7}$\\,\\smy\\ and $r_{out}=4\\times10^{4}$\\,\\rg, equation~\\ref{eq:fkr1} indicates that $L_{5500}^{(td)} \\la 10^{37}$\\,\\ergs\\,\\AA$^{-1}$ for $\\rtrans\\ga10^{2}$\\,\\rg. {\\it At the accretion rates required by the ADAF models, the outer thin disk must extend to within $\\sim10^{2}$\\,\\rg\\ in order to generate sufficient optical luminosity. This is additional, independent evidence that $\\rtrans\\la10^{2}$\\,\\rg\\ in NGC 4258.} The masers in NGC 4258 extend from 0.13 pc to 0.26 pc, and the cause of the observed inner edge at $4\\times10^{4}$\\,\\rg\\ is uncertain. The present observations effectively disqualify one of the most promising theories: that the inner edge of the masing zone corresponds to \\rtrans\\ and to the inner edge of the outer cool disk itself. Even if NGC 4258 is powered by an ADAF, the lack of 22-GHz emission requires that the outer thin disk extends to well within $4\\times10^{4}$\\,\\rg, and another explanation is needed to account for the lack of masers in this region. We note that Neufeld \\& Maloney (1995) argue that a proposed flattening of the disk at small radii would reduce the disk's exposure to central, hard X-rays, and may account for the lack of maser emission within 0.13\\,pc. There are a great many low-luminosity accreting systems, encompassing a wide variety of types and spanning many decades in central mass, that are apparently sufficiently sub-Eddington to be plausible ADAF candidates. However, for most of these systems, a standard Shakura-Sunyaev thin disk is also a viable solution. Thus, while there is compelling evidence for the existence of ADAFs in some systems, it remains unclear to what extent the optically thin, two-temperature ADAF is in general applicable. Unfortunately, a physically motivated explanation for how ADAFs are triggered in the first place remains elusive, and, until this `on-switch' is discovered, questions concerning the universality of ADAFs must necessarily be addressed empirically. NGC 4258 is a useful place to explore this issue, since it is highly sub-Eddington, and because the nuclear maser and VLBI together provide access to a host of usually obscure parameters. While the present 22-GHz non-detection does not rule out the presence of an ADAF in NGC 4258, it does begin to place interesting constraints on the geometry of any proposed ADAF." + }, + "9802/astro-ph9802114_arXiv.txt": { + "abstract": "We study the formation of molecular hydrogen in cooling gas behind shocks produced during the blow-away process thought to occur in the first collapsed, luminous (Pop~III) objects in the early universe. We find that for a wide range of physical parameters the $H_2$ fraction is $f \\approx 6 \\times 10^{-3}$. The $H_2$ mass produced in such explosions can exceed the amount of relic $H_2$ destroyed inside the photodissociation region surrounding a given Pop~III. We conclude that these first objects, differently from the suggestion of Haiman \\etal 1997, might have a net {\\it positive} feedback on subsequent galactic formation. We discuss the effects of radiation and the implications of our results for the soft-UV background. ", + "introduction": "Current models of cosmic structure formation based on CDM scenarios predict that the first collapsed, luminous (hereafter Pop~III) objects should form at redshift $z\\approx 30$ and have a total mass $M \\approx 10^6 M_\\odot$ or baryonic mass $ M_b \\approx 10^5 M_\\odot$ (Couchman \\& Rees 1986, Haiman \\etal 1997 [HRL], Tegmark \\etal 1997). This conclusion is reached by requiring that the cooling time, $t_c$, of the gas is shorter than the Hubble time, $t_H$, at the formation epoch. In a plasma of primordial composition the only efficient coolant in the temperature range $T\\le 10^4$~K, the typical virial temperature of Pop~III dark matter halos, is represented by H$_2$ molecules whose abundance increases from its initial post-recombination relic value to higher values during the halo collapse phase. It is therefore crucial to determine the cosmic evolution of such species in the early universe to clarify if small structures can continue to collapse according to the postulated hierarchical structure growth or if, lacking a cooling source, the mass build-up sequence comes to a temporary halt. The appearance of Pop~III objects is now thought to cause a partial destruction of the available molecular hydrogen either in the intergalactic medium (IGM) and/or in collapsing structures; the result is a negative feedback on galaxy formation. This effect has been pointed out by HRL and it works as follows. As stars form in the very first generation of objects, the emitted photons in the energy band 11.2-13.6 eV are able to penetrate the gas and photodissociate H$_2$ molecules both in the IGM and in the nearest collapsing structures, if they can propagate that far from their source. This negative feedback and its possible limitations are discussed by Ciardi, Abel \\& Ferrara (1998, CAF). Here we propose and investigate instead a possible {\\it positive} feedback based on supernova (SN) explosions, which, under many aspects, is reminiscent of a scaled version of the explosive galaxy formation scenario introduced by Ostriker \\& Cowie (1981) and put forward by many others. Pop~IIIs are very fragile due to their low mass and shallow gravitational potential: only a few SNe are sufficient to blow-away (Ferrara 1998) their baryonic content and drive an expanding blastwave into the IGM, which eventually becomes radiative and allows the swept gas to cool in a dense shell. The cooling transient, as we will see, is characterized by a strong nonequilibrium condition in which recombination lags behind the temperature decrease. As already pointed out by Shapiro \\& Kang (1987) and Kang \\& Shapiro (1992), this is a favorable condition for H$_2$ formation. Our conclusions are that the amount of molecular hydrogen thus formed can exceed the one destroyed via photodissociation, yielding a net increase of the H$_2$ in the universe. As a consequence, not only the galaxy formation process is not halted, but instead it is favored by the effects of Pop~III formation. Sec. 2 describes the main properties of multi-SN explosions in the early universe; Sec. 3 is devoted to the calculation of H$_2$ formation in their cooling shell. In Sec. 4 we compare the magnitude of positive and negative feedbacks; some discussion of the results is given in Sec. 5. ", + "conclusions": "The results obtained in the previous Section could be modified by the fact that so far we have neglected the effects of the Pop~III stellar cluster radiation impinging on the cooled shell. In fact, even after all the SNe have exploded blowing-away the gas, the coeval low mass stars will continue to produce some residual flux in the energy range $0.755-13.6$~eV, relevant to the $H_2$ formation network. Two processes might have some effect in this context: (i) $H^-$ photoionization, and (ii) $H_2$ two-step photodissociation. The first process occurs at a rate $k_{23} \\approx 6\\times 10^{-17} S_{LW}/R_s^2(t_c)$~s$^{-1}$ per $H^-$ atom, or $2.4\\times 10^{-10} \\beta M_6^{3/5}$~s$^{-1}$ using eqs. \\ref{Rsc}, \\ref{S} and reference values for the remaining parameters. The most effective $H^{-}$ destruction process among the ones included is $H_2$ formation, whose rate in the same case is $\\sim 1.4 \\times 10^{-9}$~s$^{-1}$. Thus, photoionization is negligible, and we have checked this conclusion numerically. We can calculate the effects of two-step photodissociation on the shell as follows. The shell $H_2$ column density is $N_{H_2} \\simeq f M_s/m_p R_s(t_c)^2 = 1.5\\times 10^{17} (1+z)_{30}^{11/10} M_6^{1/5}$~cm$^{-2}$ (we have assumed that all the swept IGM mass is in the cool shell). Since the critical column density for self-shielding is $N_{H_2}^{crit}=10^{14}$~cm$^{-2}$ (Draine \\& Bertoldi 1997), the shell is optically thick to dissociating radiation. Nevertheless, the flux in the LW bands, $J_d =S_{LW} h_{_P}/ 4\\pi R_s^2$, where $h_{_P}$ is the Planck constant, will induce a dissociation front in the shell; $J_d$ is very likely dominating on a possible background radiation. The propagation speed will be $v_d \\simeq \\xi \\lambda_d/t_d$, where $\\lambda_d= N_{H_2}^{crit}/n_{H_2}$ is the LW photon mean free path, and $\\xi \\approx 3-5$ is a constant obtained by comparison with numerical results. Thus, the time required to dissociate the entire shell is $\\simeq N_{H_2}/n_{H_2} v_d = 3\\times 10^8 (J_{d,21} \\xi)^{-1}N_{H_2,17}$~yr, where $J_{d,21}=J_{d}/ 10^{-21} {\\rm ergs~cm}^2{\\rm s}^{-1}{\\rm Hz}^{-1} {\\rm sr}^{-1}$, and $N_{H_2,17}= N_{H_2}/10^{17} {\\rm cm}^{-2}$. Since theoretical work (Bruzual \\& Charlot 1993) suggests that $\\beta$ drops to $\\approx 3\\times 10^{-3} \\ll 1$ after all the massive stars (the main contributors to the UV flux) have died, dissociating the entire shell will require a very long time, certainly larger than the Hubble time. Finally, Kang \\& Shapiro (1992) pointed out that any external {\\it ionizing} radiation field (which could be provided by the hot postshock gas and/or by the residual intermediate mass stars in the cluster) tends to increase the final value of $f$, although introducing a time delay due to a temperature plateau in the evolution characteristic of photoionization heating. Thus, our estimate seems to provide a robust lower limit to the molecular hydrogen abundance. How can this $H_2$ enhanced gas be used for galaxy formation ? First, regions in which Pop~IIIs are clustered enough for their shells to interact might become sites of active star/galaxy formation. Shell interactions appear to be necessary to make them gravitational unstable, since the temperature required for isolated fragmentation (Couchman \\& Rees 1986) is well below the values ($ \\simgt 100$~K) found here. Next, these intergalactic shells can be accreted by neighbor objects, possibly increasing their $H_2$ abundance above the threshold $f\\simgt 5\\times 10^{-4}$ required for efficient cooling (Tegmark \\etal 1997). We conclude that Pop~III objects can produce regions of considerably high molecular hydrogen abundance due to multi-SN shocks propagating in the IGM. We have also seen that the $H_2$ thus produced can exceed the amount of relic $H_2$ destroyed inside the photodissociation region surrounding a given Pop~III. This occurrence suggests that these first objects might have a positive feedback on subsequent galactic formation. Such conclusion is certainly valid on a local scale, as defined by the radii of the two influence spheres, $R_s$ and $R_d$, of a single object. However, our definition of $R_d$ (see Sec. 4) does not exclude that some photodissociating flux is present outside $R_d$. If many of these objects are present in the universe at the same epoch, they might contribute to an early soft-UV background which, illuminating an isolated, collapsing cloud, could photodissociate and depress its $H_2$ content; this possibility has been discussed by HRL. However, CAF have demonstrated that, once the {\\it surviving relic IGM fraction between photodissociated spheres} is taken into account, the universe opacity to LW photons before reionization becomes of order unity on scales larger than $R_d$, but smaller than the typical interdistance between Pop~IIIs at $z=20-30$. This implies that the soft-UV background is very weak due to intergalactic absorption, and, as a consequence, the proposed negative feedback considerably inhibited. Furthermore, most of the dissociating flux is absorbed by the thick $H_2$ shells discussed above. Given that in general $R_d > R_i$, where $R_i$ is the HII region radius of a Pop~III (CAF), clearing of intergalactic $H_2$ by photodissociated sphere overlapping will occur before complete reionization. It is not clear at this stage if that event occurs before the typical mass scale of the collapsing objects is such that cooling can proceed via Ly$\\alpha$ only, thus considerably weakening the arguments in favor of a temporary halt of galaxy formation at high redshift. \\vskip 2truecm We acknowledge T. Abel, B. Ciardi, D. Cox and G. Field for discussions; E. Corbelli, Z. Haiman, A. Loeb, F. Palla, O. Pantano, P. Pietrini, and Y. Shchekinov for comments; thanks are due to L. Pozzetti for help with the spectrophotometric code. \\vskip 2truecm" + }, + "9802/astro-ph9802028_arXiv.txt": { + "abstract": "\\noindent Phase transitions in neutron stars due to formation of quark matter, kaon condensates, etc. are discussed with particular attention to the order of these transitions. Observational consequences of phase transitions in pulsar angular velocities are examined. ", + "introduction": "The physical state of matter in the interiors of neutron stars at densities above a few times normal nuclear matter densities is essentially unknown. Interesting phase transitions in nuclear matter to quark matter \\cite{Glendenning,HPS}, kaon \\cite{Kaplan,kaon} or pion condensates \\cite{pion,vijay}, neutron and proton superfluidity \\cite{oeystein}, hyperonic matter, crystalline nuclear matter \\cite{pion}, magnetized matter, etc., have been considered. We discuss how these phase transitions may exist in a mixed phase, the structures formed and in particular the order of the transition. Observational consequences are discussed. ", + "conclusions": "We have discussed various possible phase transitions in neutron stars and have argued that we expect several first order order phase transitions to occur when the topological structure of the mixed phase change in the inner crust, nuclear and quark matter mixed phase or Kaon condensates. If a first order phase transitions is present at central densities of neutron stars, it will show up in moment of inertia and consequently also in angular velocities in a characteristic way. For example, the braking index diverges as $n(\\Omega)\\sim c_2/\\sqrt{1-\\Omega^2/\\Omega_0^2}$." + }, + "9802/astro-ph9802245_arXiv.txt": { + "abstract": "We present a thorough quantitative analysis of the evolution of the colour-magnitude relation for early-type galaxies in 17 distant clusters with redshifts $0.31 < z < 1.27$ using the Kodama \\& Arimoto (1997) evolutionary model for elliptical galaxies. The model is calibrated to reproduce the colour-magnitude relation for Coma ellipticals at $z\\sim 0$ and gives the evolution of the slope and zero-point as a function of redshift. We find no significant differences between the colour-magnitude relations of the clusters in our sample. The slopes can be reproduced by a single model sequence in which all elliptical galaxies are assumed to be equally old (the maximum age difference allowed for the brightest 3 magnitudes is only 1~Gyr) and to have mean stellar metallicities which vary as a function of galaxy luminosity. The zero-points of the colour-magnitude relations constrain the epoch of major star formation in early-type galaxies to $z_f>2$--$4$. This study provides two important constraints for any model of the formation of rich clusters: the uniformity of the ages of the stellar populations in the early-type galaxies and the universality of the metallicity sequence of these galaxies as a function of galaxy mass. ", + "introduction": "Visvanathan \\& Sandage (1977) first noted from their data on 9 nearby clusters that the more luminous early-type galaxies tended to have redder colours. Bower, Lucey, \\& Ellis (1992; hereafter BLE92) later studied this {\\it colour-magnitude} ({\\it C-M}) relation in detail using high precision photometry of early-type galaxies in the Virgo and Coma clusters and found very little scatter about the mean {\\it C-M\\/} relation. Thus, there appeared to be a marked homogeneity in the present-day early-type cluster galaxy population. Recently, Kodama \\& Arimoto (1997; hereafter KA97) compared the predictions of their evolutionary models of elliptical galaxies with the {\\it C-M\\/} diagrams for early-type galaxies in two distant clusters, Abell 2390 at $z=0.228$ and Abell 851 at $z=0.407$, and showed that the origin of the {\\it C-M\\/} relation was primarily due to the mean stellar metallicity of the early-types being a function of total magnitude. In this paper we present a more thorough study of the evolution of the {\\it C-M\\/} relation by using 17 clusters at cosmological distances to detect the colour evolution of cluster early-type galaxies directly and then comparing the results with the models. This will enable us to put even stronger constraints on the formation of early-type galaxies in clusters. The photometric evolution of early-type galaxies in $z<1$ clusters was previously examined by several authors (e.g., Ellis et al.\\ 1985; Couch, Shanks \\& Pence\\ 1985; Arag\\'on-Salamanca et al.\\ 1993). Arag\\'on-Salamanca et al.\\ (1993) traced the evolution of the `red envelope' from the optical-near-infrared colour distribution of galaxies in 10 clusters with redshifts $0.5 < z < 0.9$ and concluded that the detected evolution was consistent with the passive ageing of stellar populations formed before $z\\simeq 2$. A potential problem with these studies is that the early-type galaxies had to be selected by their spectral energy distributions and colours, thereby increasing the possibility of contamination from other galaxy types. The advent of the Hubble Space Telescope {\\it (HST)}, with its spectacular high-resolution images, has made the morphological classification of distant galaxies possible. Furthermore, thanks to the {\\it HST}, large ground-based telescopes, and the high quality of astronomical detectors, much fainter limiting magnitudes can now be reached. An additional advantage of {\\it HST} photometry comes from the reduction of the sky background at longer wavelengths compared to that achieved with ground-based observations (e.g. the background in space is 8 times fainter in the $I$-band), thereby reducing the photometric errors at faint limits significantly (Ellis et al.\\ 1997; hereafter E97). These improvements now make it possible to examine the evolution of early-type galaxies over a much wider luminosity and redshift range. The {\\it C-M\\/} relation is a powerful tool for quantifying the colour evolution of early-type galaxies as a function of redshift since it can be observed up to very high redshifts. Dickinson (1996) showed that the {\\it C-M\\/} relation is already recognizable at $z\\simeq 1.2$. The zero-point of the {\\it C-M\\/} relation at high redshift provides direct information on the properties of early-type galaxies in general. On the other hand, the slope of the {\\it C-M\\/} relation has information on differential properties of early-type galaxies as a function of luminosity or galaxy size. Stanford, Eisenhardt, \\& Dickinson (1995; 1998, hereafter SED98) presented {\\it C-M\\/} relations for morphologically selected early-type galaxies in 19 clusters out to $z=0.9$. When they compared their optical-near-infrared {\\it C-M\\/} relations to that of Coma, they found no significant change in either the slope or the scatter as a function of redshift. They did, however, observe a progressive blueing with redshift of the average colour in a manner consistent with the passive evolution of an old stellar population formed in a burst at an early cosmic epoch. E97 also analyzed the {\\it C-M\\/} relations of three distant clusters at $z\\simeq 0.54$ (Cl~0016$+$16, Cl~ 0054$-$27, Cl~0412$-$65), in this case using {\\it HST} photometry. They found that the dispersion about the relations was quite small even in such distant clusters, thereby requiring a high formation redshift for the stars in early-type galaxies, such as $z\\simeq3$. E97 also did not detect a significant change in the slopes of the relations in the rest-frame $U-V$ when compared with Coma. The zero-points show modest colour evolution in agreement with earlier studies (Ellis et al. 1985; Couch et al. 1985; Arag\\'on-Salamanca et al. 1993). All of these results support a formation picture of early-type galaxies in which the bulk of the star formation is completed at high redshift with little star formation occurring in the recent past. Thus, the {\\it C-M\\/} relation originates at high redshift as a metallicity effect. Recently, however, Kauffmann \\& Charlot (1997) have argued that the {\\it C-M\\/} relation can alternatively be interpreted in the context of the hierarchical models of galaxy formation once chemical enrichment is taken into account. In their analysis, the slope of the {\\it C-M\\/} relation is maintained because large ellipticals form primarily from large metal-rich progenitors. They find that it remains nearly constant up to $z\\simeq 1$. Furthermore, they are able to explain the tightness of the {\\it C-M\\/} relation, despite frequent galaxy merging, by suggesting that we are biasing ourselves to the selection of only old galaxies by studying rich clusters at high redshift. According to their models, these objects are not the progenitors of present-day clusters. It is not yet clear whether this type of model can give an universal {\\it C-M\\/} relation for different clusters, since merging histories are likely to have varied from cluster to cluster. Our motivations to conduct further analyses on the evolution of the {\\it C-M\\/} relation for cluster early-type galaxies are twofold: (1)~to investigate the universality of the {\\it C-M\\/} relation in clusters and (2)~to put stronger constraints on the formation of early-type galaxies by using the more realistic elliptical galaxy model of KA97, which takes into account chemical evolution, rather than an ad-hoc single burst model with fixed metallicity ($Z_\\odot$) adopted by most of the previous analyses. KA97 have constructed an evolutionary model of elliptical galaxies which uses a population synthesis technique based on a monolithic collapse picture of galaxy formation. The marked advantage of the KA97 model is that it allows us to analyze both the slope of the {\\it C-M\\/} relation and the zero-point, since it includes the effects of metallicity. Moreover, since the model is calibrated with the empirical colours of ellipticals in Coma, it can be directly and reliably compared to the observational data in distant clusters. In this paper, we investigate the evolution of the {\\it C-M\\/} relation of 17 distant clusters using the KA97 evolutionary models. We have accumulated the photometric data from the literature, most of which were obtained with the {\\it HST} and large, ground-based telescopes ($\\sim 2$--$4\\,$m). Although some of the {\\it HST} data suffer from zero-point uncertainties due to the paucity of observed standard stars (see below), their random photometric errors are very small down to $\\sim 3\\,$mag from the brightest end of the {\\it C-M\\/} relation (typically $0.01-0.07\\,$mag). Thus, we are able to conduct a reliable analysis of, for example, the slope of the {\\it C-M\\/} relation based on the relative photometry within each cluster. To constrain the formation epoch of early-type galaxies in clusters, we also conduct a zero-point analysis using only those clusters which have been calibrated with ground-based photometry or those which are at sufficiently high redshift that the evolutionary changes are large enough to put some constraint on the formation epoch, despite relatively large zero-point uncertainties. The cosmological parameters we have chosen to use throughout this paper are $H_0=50\\,$km$\\,$s$^{-1}\\,$Mpc$^{-1}$ ($h\\equiv H_0/100 = 0.5$) and $q_0=0.5$, without a lambda term, unless otherwise stated. The structure of this paper is as follows. In \\S2 we compile the observational data and define the observed {\\it C-M\\/} relation for each cluster. In \\S3 we give a brief description of the evolutionary models of KA97. We compare the models and the observations in \\S4, and in \\S5 we give a discussion of the results and our conclusions. ", + "conclusions": "We have shown that the {\\it C-M\\/} relations of early type galaxies in clusters with $0.311$ even with moderately accurate photometry (errors$\\sim0.1\\,$mag). Differences of $\\sim1\\,$Gyr are easily measurable at such high redshifts. Therefore, photometry of early-type galaxies in clusters beyond $z\\sim1$ can accurately determine the formation epoch of these galaxies. Furthermore, the slope of the {\\it C-M\\/} relation at redshifts beyond $z\\sim1$ is so sensitive to age difference (as suggested by the long-dashed lines in the bottom right panel of Fig.~\\ref{fig:slope}) that it will allow us to investigate systematic difference of mean stellar age as small as $\\simeq$1~Gyr as a function of galactic mass, if they are present. We expect slope changes to be especially prominent when we approach the star formation phase of the elliptical galaxies. We have not discussed the colour scatter around the {\\it C-M\\/} relation in this paper, but SED98 pointed out that the scatter dose not increase with redshift out to $z\\sim1$. It would be extremely important to determine when the {\\it C-M\\/} relation breaks down: i.e. when its scatter becomes very large. This would indicate the very time when the ellipticals are forming. Unfortunately, only a handfull of clusters have been found beyond $z\\sim1$. Searching for such distant clusters and pushing the {\\it C-M\\/} relation analysis towards higher redshifts must be a promising strategy to determine the formation epoch of cluster ellipticals globally and as a function of their mass. The advent of large format near-infrared arrays provides the means to carry out that search efficiently. We finish with a word of caution. Our study is based on the monolithic model for the formation of early-type galaxies. The data we have analyzed here is fully consistent with this model, but it might not be the only solution. In particular, an alternative scenario set in the context of the hierarchical merging model of galaxy formation (e.g. Kauffmann \\& Charlot 1997) could give results that are also broadly consistent with the available data, although it has not been fully confirmed yet. A detailed comparison of the observed properties of the {\\it C-M\\/} relation in distant clusters with the predictions of hierarchical models that properly take chemical evolution into account is clearly needed. It is especially important to test whether this type of model can give the universal {\\it C-M\\/} relation for different clusters locally and over a wide range of redshift, since merging histories are likely to have varied from cluster to cluster. Moreover it is not yet confirmed that larger ellipticals form always from larger spirals already enriched in metal. In any case, however, if we define the formation epoch of the early-type galaxies as the time when the bulk of their stars formed, the properties of the {\\it C-M\\/} relation place that epoch well beyond $z=2$, independent of which formation picture is correct." + }, + "9802/astro-ph9802073_arXiv.txt": { + "abstract": "The structure of steady plane-parallel radiative shock waves propagating through the hydrogen gas undergoing partial ionization and excitation of bound ato\\-mic states is investigated in terms of the self-consistent solution of the equations of fluid dynamics, radiation transfer and atomic kinetics. The shock wave model is represented by a flat finite slab with no incoming radiation from external sources at both its boundaries. The self-consistent solution is obtained using the global iteration procedure each step of which involves (1) integration of the fluid dynamics and rate equations for the preshock and postshock regions, consecutively, both solutions being fitted by the Rankine-Hugoniot relations at the discontinuous jump; (2) solution of the radiation transfer equation for the whole slab. The global iteration procedure is shown to converge to the stable solution which allows for the strong coupling of the gas flow and the radiation field produced by this flow. Application of the method is demonstrated for the shock waves with upstream velocities of $15~\\kms\\le U_1\\le 60~\\kms$ (i.e. with upstream Mach numbers $2.3\\le M_1\\le 9.3$) and the hydrogen gas of unperturbed temperature $T=3000\\K$ and density $\\rho = 10^{-10}~\\gmcs$. ", + "introduction": "Radiative shock waves belong to the conspicuous phenomena demonstrating the tight interplay between hydrodynamic motions and the radiation field. The role of this interplay is strongest in the low gas density flows, so the shocks are of tremendous importance in astrophysics. They are observed in a wide variety of astrophysical phenomena: nova and supernova explosions, bright filaments in old supernova remnants, accretion flows in protostellar clouds. Shock waves are detected also in atmospheres of radially pulsating variables such as Cepheids, RR~Lyr, W~Vir, RV~Tau and Mira type stars. Periodic shocks propagating through pulsating atmospheres lead to the distention of outer atmospheric layers and to the mass loss. Importance of radiative shock waves attracted attention of many authors but nevertheless the shock properties are explained quite well still qualitatively (see, for example, Zel'dovich \\& Raizer 1966; Skalafuris 1968; Mihalas \\& Mihalas 1984; Liberman \\& Velikovich 1986). The principal difficulty in obtaining the correct quantitative description of the shock wave structure is that the model has to allow for the strong coupling between the gas flow and the radiation field, both them being characterized by substantial departures from LTE. Solution of this problem encounters serious difficulties, so that in immensely numerous studies available at present in the literature the authors used various assumptions and simplifications (e.g. local thermodynamic equilibrium, treatment of the radiation transfer in diffusion approximation, neglecting the opacity in the Balmer continuum etc.). Many of these assumptions were found later inadequate or leading to uncertain conclusions. For instance, Kogure (1962), Sachdev (1968) and Hill (1972) used the LTE approximation which does not hold as emphasized later Narita (1973). Considering the hydrogen gas, Whitney \\& Skalafuris~(1963) relaxed this assumption but incorrectly assumed that the postshock region is transparent for all hydrogen continua. Finally, Narita (1973) took into account the opacity in the both Lyman and Balmer continua. The most elaborate numerical modelling of radiative shock waves based on the self-consistent solution of the equations of fluid dynamics, radiative transfer and atomic level populations was done by Klein et al. (1976, 1978). However, the coarse zonning $(\\sim 10^6~\\cm)$ did not allow to authors to consider the detailed structure of the shock front including the radiative precursor and the thermalization zone where the electron temperature gradually equalizes with temperature of heavy particles. Nevertheless, this approach was found to be enough for consideration of shock dynamics in atmospheres of A-type stars because the radiative precursor is not so important due to the high temperature of the unperturbed gas. Because astrophysical shocks in stellar atmospheres propagate through the partially ionized hydrogen gas, a substantial fraction of photons produced within the wake are absorbed in the radiative precursor. As was shown by Gillet \\& Lafon (1984, 1990) the structure of the radiative precursor is complex and should be treated with same degree of approximation as the postshock region. In their studies Gillet \\& Lafon (1984, 1990) treated the radiative transfer as an initial value problem which was solved using the shooting method. The principal obstacle in such an approach is that the transfer equation possesses a singularity in the postshock region (Gillet et al. 1989). Indeed, application of the eigenvalue methods for solution of the nongrey transfer problem is affected by exponentially growing errors (Mihalas 1978). One of the first attempts to obtain the self-consistent solution for the shock wave structure was undertaken by Nelson \\& Goulard (1969) and Nelson (1973). They considered the shock waves propagating through the argon-like gas with upstream Mach numbers of $M_1=18$ and $M_1=24$. The continuity, momentum and energy equations were written in the integral representation whereas the radiation transfer was treated in the simplified formulation. The studies of radiative shock waves in helium and nitrogen done by Clarke \\& Ferrari (1965), Farnsworth \\& Clarke (1971) and Foley \\& Clarke (1973) seem to be the best among known in the literature. The authors emphasized the crucial role of the radiation transfer treatment and employed the formal solution of the transfer equation. The self-consistent shock wave models were obtained in these studies with iteration procedure. Unfortunately, there is a problem of exponential factors when the formal solution is applied for optically thick layers. In this paper we present a new approach based on the iterative solution of the equations of fluid dynamics, the rate equations and the radiation transfer equation. The momentum equation, the energy and rate equations are written in the form of ordinary differential equations. These equations are stiff and such a representation is most appropriate from the point of view of stability and small truncation errors. The radiation transfer is treated as a two-point boundary value problem. This allows us to obtain the stable solution of the transfer equation for the whole spectral range including both the opaque Lyman continuum and the more transparent higher order continua. The method of global iterations takes into account the coupling between the gas flow and the radiation field, so that the structure of the radiative shock wave is considered in terms of the self-consistent model. In the framework of this first approach, only devoted to provide a new technique for obtaining the self-consistent solution, we consider the structure of steady, plane-parallel shock waves propagating through an infinite, isotropic, pure hydrogen plasma. The steady assumption is correct, for example, to a good degree of approximation in most applications to stellar atmospheres. Indeed, the time required for the gas flow to cross the shock wake with typical thickness of 10$^{4}$-10$^{7}$~cm is much less than the characteristic time during of which the shock wave energy appreciably decreases. In pulsating stars, the radiative lifetime of a shock is between a few hours and a few months, which, consequently, is much larger than the 10$^{-2}$-10~s of the gas flow time to cross the shock wake. Accuracy of the plane-parallel approximation follows from the very small width of the shock wave in comparison with stellar radius. After describing the shock wave model (Sect.~2) we derive the system of ordinary differential equations (Sect.~3). The radiation transfer equation is solved for the whole shock wave model using the Feautrier technique (Sect.~4). In Sect.~5 we show that the global iteration procedure comprising the initial value problem for ordinary differential equations and the two-point boundary value problem for radiation transfer converges to the self-consistent solution. Results of calculations demonstrating the applicability of the method are given in Sect.~6. Finally, in Sect.~7, we give some concluding remarks and discuss the future aspects of the problem. ", + "conclusions": "The primary goal of the work reported in this paper was to obtain the self-consistent solution of the equations of fluid dynamics, rate equations and radiation transfer equation for the structure of the steady shock wave. The procedure of global iterations described in the present paper in general resembles the compute of stellar atmosphere models. Indeed, like in stellar atmosphere calculations the shock wave model takes into account the coupling between the gas material and radiation field. The self-consistent model is obtained with iteration procedure comprising the solution of the radiation transfer equation and integration of the mass, momentum and energy conservation equations written in the form of the ordinary differential equations. Each cycle of iterations gives, in general, improved characteristics of the gas and radiation field. At the same time, the problem of the shock wave structure compared to that of stellar atmosphere models contains a number of serious complications. First, atomic level populations are not only in strong departures from LTE but are also in significant departures from statistical equilibrium. Second, unlike the stellar atmospheres, where the divergence of radiative flux is $\\df=0$ (the condition of radiative equilibrium), in shock waves the part of the energy of hydrodynamic flow is transformed into radiation and the radiative equilibrium is established only far away from the discontinuous jump. Furthemore, in stellar atmosphere models the total radiative flux is given as one of the boundary conditions, whereas in the shock wave model the emerging flux is obtained from the solution of the problem. The small optical depth increments in hydrogen continua of order $l\\ge 3$ lead to the losses of the machine accuracy when the Feautrier technique is applied. The problem is so serious, that even the improved method by Rybicki \\& Hummer (1991) sometimes fails. Third, the rate equations are stiff and need the special treatment in their solution. In particular, the convergence of global iterations depends on the tolerance parameter determining the maximum error permitted during the integration. The present paper is confined by consideration of the two-level atomic model, so that the radiation transfer is treated for the Lyman and Balmer continua, only. This approximation seems to be insufficient for the shock wave problem because the occupation numbers of levels $l\\ge 2$ obviously deviate from LTE and the perceptible fraction of radiation is transported at frequencies lower than the Balmer edge frequency. Nevertheless, notwithstanding such a restriction, there is a qualitative agreement of our results with those obtained earlier by other authors. For example, according to calculations of Gillet \\& Lafon (1990) the electron temperature just ahead the discontinuous jump is $\\Te\\approx 14000\\K$ for the upstream velocity of $U_1=80~\\kms$. Although in the present study the highest upstream velocity was $U_1=60~\\kms$, a very approximate comparison can be done with fitting formula (\\ref{prec_te}) which gives the same electron temperature for the upstream velocity $U_1\\approx 75~\\kms$. Thus, more detailed calculations are needed and in the forthcoming paper we are going to present the grid of the shock wave models computed for the larger number of hydrogen atomic le\\-vels and wider range of upstream velocities, the models of the present study being used as initial approximation for more correct shock wave models. More realistic models, however, should take into account not only bound-free terms but also bound-bound terms in the rate equations and the radiation transfer problem should be solved for the both continuum and spectral line radiation. This is the perspective for the near future. It is certainly one of the most basic. Indeed, preceding shock studies show that radiative processes, which determine the wake cooling, have a strong influence on the resulting shock structure. Because in the model of this paper we consider a pure H-plasma without H$^{-}$, and only include the bound-free photo- and collisional processes of H atoms, we expect that the absence of some predominant coolants such as neutral and singly ionized metal atoms might appreciably underestimate the radiative cooling rate of the gas. The importance of radiative heating and cooling rates in shocked circumstellar envelopes have been recently investigated (Woitke et al. 1996) but only a few transitions of the numerous metal lines were considered. At present such a basic study seems to be beyond our immediate abilities and is out of the scope of our paper which was to investigate the possibility of obtaining a self-consistent solution of the structure of radiative shock waves in dense atmospheric gas." + }, + "9802/hep-ph9802220_arXiv.txt": { + "abstract": "In view of the ongoing galactic (or cosmic) axion detection experiments, we compare the axion-photon-photon coupling $c_{a\\gamma\\gamma}$'s for various invisible (or very light) axion models. ", + "introduction": " ", + "conclusions": "" + }, + "9802/astro-ph9802186_arXiv.txt": { + "abstract": "Derived from first physical principles, a few simple rules are presented that can help in both the planning and interpretation of CCD and IR-array camera observations of resolvable stellar populations. These rules concern the overall size of the population sampled by a frame as measured by its total luminosity, and allow to estimate the number of stars (in all evolutionary stages) that are included in the frame. The total luminosity sampled by each pixel (or resolution element) allows instead to estimate to which depth meaningful stellar photometry can be safely attempted, and below which crowding makes it impossible. Simple relations give also the number of pixels (resolution elements) in the frame that will contain an unresolved blend of two stars of any kind. It is shown that the number of such blends increases quadratically with both the surface brightness of the target, as well as with the angular size of the pixel (or resolution element). A series of examples are presented illustrating how the rules are practically used in concrete observational situations. Application of these tools to existing photometric data for the inner parts of the bulge of M31, M32 and NGC 147 indicates that no solid evidence has yet emerged for the presence of a significant intermediate age population in these objects. ", + "introduction": "The use of sophisticated software packages for the photometry of stars in crowded fields is now a widespread activity. Routines take a CCD frame, and automatically flatfield it, subtract dark, bias and sky, remove cosmic ray events, restore bad pixels, and finally deliver a catalog of star magnitudes and positions. To some extent, what is left to the astronomer is to plot color-magnitude diagrams, sort out the astrophysical inperpretation of them, and finally write the paper. The use of these packages has thus resulted in a tremendous progress in the study of resolved stellar populations, including those in galactic and extragalactic globular clusters, the Magellanic Clouds, the bulges of the Milky Way and M31, the dwarf satellites of M31, and several irregular galaxies in the outskirts of the Local Group and beyond. Such progress has been possible thanks to the large size of the analyzed samples of stellar populations, coupled with a superior photometric accuracy. No doubt such achievements would have been impossible with the old, traditional method of direct optical inspection of each stellar image. However, relying entirely on the automatism of the procedures can also produce nonsense, and occasionally it really did so. In this paper it is argued that a straighforward application of some simple, yet basic physical concepts can greatly help the astronomer's work on stellar populations. Some simple tools will be presented that allow one to easily get the required numbers out for both the planning for a most efficient use of the available telescope facilities, as well as for achieving a deeper understanding of the scientific meaning of the data once they have been obtained. The paper is organized as follows. Section 2 presents a few basic conceptual tools whose application to CCD photometry of stellar populations is then illustrated in the following sections. Section 3 deals with the overall size of a sampled stellar population, hence with the stellar evolutionary phases that can be investigated given such size. Section 3 also deals with the size of the stellar population that is sampled by each resolution element, hence with the depths (limiting magnitude) that can be achieved before crowding effects hamper meaningful stellar photometry. In Section 4 several examples are presented to illustrate the practical use of such tools, and Section 5 summarizes the main conclusions of this paper. ", + "conclusions": "Simple tools have been presented that may help the astronomer wishing to extract as much sound science as possible from CCD/IR array data on stellar populations. The main conclusions of this paper can be summarized as follows. (1) While planning the CCD/IR array observations of stellar populations in globular clusters, in the Galactic Bulge, or in nearby resolvable galaxies the total luminosity sampled by the frame (in $\\lsun$ units) should be estimated from the known surface brightness distribution of the target objects. (2) On the basis of this ``frame luminosity\", the luminosity sampled by each pixel of the camera (or by each expected resolution element) should also be evaluated. (3) On the basis of the frame luminosity the astronomer can estimate with very good accuracy the number of stars in each evolutionary phase that will be framed by the camera, thus allowing one to decide whether the sampling is statistically adequate for the specific evolutionary phases under investigation. (4) On the basis of the ``pixel\" luminosity the astronomer can estimate with very good accuracy down to which magnitude one can trust the result of individual star photometry as performed by current photometric packages. (5) Frame sampling and pixel sampling set conflicting requirements, the former asking for regions of high surface brightness to be observed in order to secure a statistically adequate sample of stars, the latter instead asking for regions of low surface brightness to be observed in order to secure reliable photometry of individual stars. The tools provided in this paper allow the astronomer to proceed very rapidly with the necessary optimization, or realize that a certain scientific goal is not reachable with the available technology. (6) After the data have been taken, the frame luminosity should be estimated directly from the frame itself, after the frame has been properly dark, bias, and sky subtracted and calibrated. This will provide the most accurate estimate of both the frame and pixel luminosity as possible. (7) A basic criterion for the limit imposed by crowding states that reliable photometry can be obtained only for those stars that are brighter than the average luminosity sampled by each pixel (resolution element). Experiments that are easy to imagine should allow one to more finely calibrate this criterion, establishing in a quantitative way -- and for any specific photometric package -- how the photometry of individual stars degrades as this limit is approached, and to which extent star counts in a given magnitude bin are contaminated by blends of stars in fainter magnitude bins. (8) Some straightforward checks should be made when photometric packages restitute exceptionally bright objects, in order to ascertain whether they are real stars rather than the result of accidental blends of fainter stars sharing the same resolution element. This second option should be carefully evaluated especially when the surface density of such objects appears to be rougly proportional to the square power of the surface brightness. (9) The {\\it artificial star} experiments as currently done by standard photometric packages may be adequate to estimate the completeness of star counts as a function of magnitude. However, they would require to repeat the experiment nearly as many times as the number of stars in the frame in order to assess the extent of the {\\it migration} of a fraction of stars towards brighter magnitudes due to blending with other stars. (10) An application of these tools to existing photometric data for the bulge of M31, M32 and NGC 147 indicates that no solid evidence has yet emerged for the presence of an intermediate age population in these objects. I wish to thank Mike Rich for extensive discussions -- over the last several years -- on the effects of crowding on the results of package photometry of Local Group galaxies. I wish also to thank Russell Cannon for a critical reading of the manuscript which resulted in an improved presentation. I am grateful to Claudia Maraston for her help in producing Figure 2. \\clearpage \\begin{deluxetable}{lrrrr} \\footnotesize \\tablecaption{Star Numbers for $\\lt = 10^{5} \\lsun$, Age = 15 Gyr, and $Z=Z_\\odot$ \\label{tbl1}} \\tablewidth{0pt} \\tablehead{ \\colhead{Evolutionary Phase} & \\colhead{$t_{j}$ (yr)} & \\colhead{$N_{j}$}} \\startdata MS & $>10^{10}$ & $1.9\\times 10^6$ \\\\ SGB & $3\\times 10^9$ & 6000 \\\\ RGB & $6\\times 10^8$ & 1200 \\\\ RGBT& $5\\times 10^6$ & 10 \\\\ HB & $10^8$ & 200 \\\\ E-AGB & $1.5\\times 10^7$ & 30 \\\\ TP-AGB & $10^6$ & 2 \\\\ LPV & $2.5\\times 10^5$ & 0.5 \\\\ P-AGB& $3\\times 10^5$ & 0.6 \\\\ PN & $2.5\\times 10^3$ & 0.005\\\\ WD & $10^9$ & 2000 \\\\ BS & $2\\times 10^9$ & 200 \\\\ BS-TP-AGB& $10^6$ & 0.06 \\\\ \\enddata \\end{deluxetable} \\clearpage \\begin{deluxetable}{lrrrrrrrr} \\footnotesize \\tablecaption{Stellar Population Sampling in the Bulge and Satellites of M31 \\label{tbl2}} \\tablewidth{0pt} \\tablehead{ \\colhead{$r$} & \\colhead{$\\mu_{\\rm B}$} & \\colhead{$\\lb$} & \\colhead{$\\lt$} & \\colhead{$N_{\\rm LPV}$} & \\colhead{$N_{\\rm RGT}$} & \\colhead{$N_{\\rm 2RGT}$}\\\\ \\colhead{} & \\colhead{mag/\\squapp} & \\colhead{$\\lsun$/\\squapp} & \\colhead{$\\lsun$/\\squapp} & \\colhead{stars/\\squapp} & \\colhead{stars/\\squapp} & \\colhead{events}} \\startdata & & & M31 Bulge & & & \\\\ \\tableline $0''-4''$ & 15 &$1.5\\times 10^6$&$3.8\\times 10^6$ &20 &410 & $-$\\\\ 2$^\\prime$& 19.8\\hfil &$1.8\\times 10^4$ &$4.5\\times 10^4$ &0.24 &4.9 & 21\\\\ 4$^\\prime$& 21\\hfil &$6.0\\times 10^3$ &$1.5\\times 10^4$ &0.06&1.6& 2 \\\\ $11'.5$ & 23 &$9.5\\times 10^2$ &$2.4\\times 10^3$ &0.01 &0.2& $-$\\\\ \\tableline & & & M32$\\;\\;$ & & & \\\\ \\tableline $5^{\\prime\\prime}$ &17\\hfil &$2.4\\times 10^5$ &$6.0\\times 10^5$ &3 &66 & $-$\\\\ $27^{\\prime\\prime}$&18\\hfil &$9.6\\times 10^4$ &$2.3\\times 10^5$ &1.2&23& $-$\\\\ $2^{\\prime}$ &21.9\\hfil &$2.6\\times 10^3$ &$6.5\\times 10^3$ &0.03&0.7 &600\\\\ $2^{\\prime}$ &24.6\\hfil &$2.2\\times 10^2$ &$5.5\\times 10^2$ &$0.003$&0.06& 6 \\\\ \\tableline & & & NGC 147$\\,\\,$ & & & \\\\ \\tableline $0''.0$& 22.55 & $1.4\\times 10^3$ & $4.3\\times 10^3$& 0.02 & 0.4 & $-$\\\\ \\enddata \\end{deluxetable} \\clearpage \\clearpage" + }, + "9802/astro-ph9802192_arXiv.txt": { + "abstract": " ", + "introduction": "The brightness distribution of gamma-ray bursts (GRBs) depends on both the intrinsic properties of the emission and the spatial distribution of sources. When analyzed globally, GRBs detected by BATSE show a clear deviation from the three-dimensional Euclidean brightness distribution (Meegan et al. 1992, 1995). However, recent analyses of special GRB subclasses showed that the deviation from the Euclidean distribution does not apply to all bursts. On the contrary, `soft' bursts with typical long durations appear to have an Euclidean brightness distribution (Pizzichini 1995, Kouveliotou 1996, Belli 1997, Pendleton et al. 1997). A commonly used spectral parameter is the {\\it burst-averaged} hardness ratio $H_{32}$, i.e., the ratio of overall fluences in the BATSE energy channels no. 2 and 3 corresponding to the 100-300~keV and 50-100~keV bands, respectively. The photon fluence-averaged ratio $H^p_{32}$, and the energy fluence-averaged ratio $H^e_{32}$ are commonly used. A first approach was to consider different GRB subclasses ranked according to their $H^e_{32}$ (e.g., Pizzichini 1995). The analysis of Kouveliotou et~al. (1996) found that GRBs with photon-fluence-averaged $H^p_{32} > 4$ and $H^p_{32} < 2$ are distributed quite differently, with the soft bursts showing an Euclidean distribution. Belli (1997) used a different phenomenological criterion to distinguish GRB classes in the hardness/duration diagram ($H^e_{32}$ vs. $T_{90}$, with $T_{90}$ the duration of 90\\% energy deposition in the 50-300~keV band) divided by the relation $H^e_{32} = 2 \\, T_{90}^{1/2}$. The long-duration subclass with $H^e_{32} > 4.5$ shows the most pronounced deviation from the Euclidean distribution. Pendleton et al. (1997, hereafter P97) used a very different approach reaching qualitatively similar conclusions. P97 select GRB {\\it individual pulses} according to the existence or lack of a high energy component above 300~keV (BATSE's channel no. 4). Two classes of GRBs with without a high energy component (HE and NHE) were identified. Very often, complex GRBs with multiple pulses show both kinds of HE and NHE pulses. The interesting conclusion of this analysis is that NHE pulses, both individual or as part of more complex GRBs, appear to have Euclidean distributions (P97). All of these analyses were carried out on a purely phenomenological level, with no direct correspondence between the selecting criterion and a physical model of emission. The main motivation of our work is to provide a physical model for the GRB observable quantities ($T_{90}, H_{32}$) to better understand the GRB brightness distribution. We use the shock synchrotron model (SSM) of GRB emission that successfully explains the great majority of GRB burst-averaged spectra (Tavani 1996a,b, 1997; hereafter T96a,b, T97). The results of our GRB analysis are obviously independent of the SSM: only the interpretation of the adopted GRB selection depends on the theoretical model. We studied GRBs of the Fourth BATSE (4B) Catalog (Paciesas et al. 1997). Three effects should be considered: (1) the intrinsic spectral evolution within bursts (typically hard-to-soft) that tends to concentrate bursts diagonally from top left to bottom right of the $H^e_{32}-T_{90}$ diagram; (2) the effect due to cosmological redshift, or any other cooling mechanism simultaneously dilating the time profile and softening the spectrum (T96a,b); (3) selection effects due to BATSE sensitivity. If the majority of GRB sources are at extragalactic distances, an overall softnening of dim and time-dilated bursts is expected as these bursts should be the most distant for no spectral cosmological evolution. The population of soft and long duration GRBs is expected be inhomogeneously distributed if BATSE's sampling radius exceeds the spatial distance of peak burst formation rate. As we show below, this is the opposite of what we find. ", + "conclusions": "" + }, + "9802/nucl-th9802003_arXiv.txt": { + "abstract": "\\noindent We calculate several ``size-like'' $^8$B observables within the microscopic three-cluster model and study their potential constraints on the zero-energy astrophysical $S_{17}(0)$ factor of the $^7{\\rm Be}(p,\\gamma){^8{\\rm B}}$ reaction. We find within our three-cluster model that a simultaneous reproduction of the experimental data for the $^8$B radius and quadrupole moment and of the $^8$B-$^8$Li Coulomb displacement energy implies $S_{17}(0)=(23-25)$ eV$\\,$b. ", + "introduction": "The $^7{\\rm Be}(p,\\gamma){^8{\\rm B}}$ reaction is currently considered to be one of the astrophysically most important nuclear reactions, as its low-energy cross section determines the high-energy solar neutrino flux \\cite{Bahcall}. Recently there has been a great deal of experimental and theoretical activities investigating this process. The low-energy cross section has been studied directly by using a radioactive $^7$Be target and a proton beam \\cite{Hammache}, in inverse kinematics by using a $^7$Be beam and a proton target \\cite{Campajola}, indirectly from the Coulomb dissociation of $^8$B \\cite{Motobayashi}, and by extracting the $^7{\\rm Be}+p$ nuclear vertex constant from the $^7{\\rm Be}(d,n){^8{\\rm B}}$ reaction \\cite{Liu}. On the theoretical side some effects of $^7$Be deformations \\cite{Nunes} and three-body dynamics \\cite{Grigorenko} have been studied, and efforts to understand the nuclear vertex constant have been made \\cite{Timofeyuk}. In Ref. \\cite{PRC} we have shown that the zero-energy cross section of the $^7{\\rm Be}(p,\\gamma){^8{\\rm B}}$ reaction scales linearly with the, unfortunately yet unknown, quadrupole moment of $^7$Be. In the present paper we extend this study and investigate the relation between the zero-energy cross section and several $^8$B ``size'' properties. In our approach we study the ground state properties of $^7$Be and $^8$B as well as the $^7{\\rm Be}(p,\\gamma){^8{\\rm B}}$ reaction cross section consistently within the microscopic eight-body $^4{\\rm He}+{^3{\\rm He}}+p$ cluster model. As this model has been discussed before we refer the reader to Refs.\\ \\cite{Csoto94,PRC} for details of the theoretical background. As customary in nuclear astrophysics we define the cross sections in terms of the astrophysical $S$ factor \\begin{equation} S(E)=\\sigma (E)E\\exp{\\Big [2\\pi\\eta (E)\\Big ]}, \\hskip 0.5cm \\eta (E)={{Z_1Z_2e^2}\\over{\\hbar v}} , \\end{equation} where $Z_1,Z_2$ are the charges of the two colliding nuclei, and $v$ is their relative velocity. At low, and in particular, at solar energies the $^7{\\rm Be}(p,\\gamma){^8{\\rm B}}$ reaction is highly peripheral, which means that only the external parts of the bound- and scattering wave functions contribute to the radiative capture cross section \\cite{Langanke}. The external wave functions are known with the exception of the asymptotic normalization constant, $\\bar c$, of the $^8$B bound state \\cite{Xu}. Consequently the energy dependence of the low-energy $S_{17}(E)$ factor is well-known (e.g. Refs. \\cite{Williams,Langanke,Csoto97}). Its absolute value, however, depends on ${\\bar c}$ and has thus to be determined experimentally. Nevertheless theoretical constraints on the asymptotic normalization constant might be quite useful. We note that $\\bar c$ depends mainly on the effective $^7$Be--$p$ interaction radius. A larger radius results in a lower Coulomb barrier, which leads to a higher tunneling probability into the external region, and hence to a higher cross section. A possible way to constrain the interaction radius is to study some key properties of the $A=7$ and 8 nuclei \\cite{ENS}. The observables that are most sensitive to the interaction radius are ``size-like'' properties, for example, quadrupole moment, radius, Coulomb displacement energy \\cite{Brown}, etc. These are the quantities which we will calculate in our microscopic cluster model and then study their effect on the astrophysical $S$ factor. ", + "conclusions": "In summary, we have adopted a microscopic $^4{\\rm He}+{^3{\\rm He}}+p$ cluster model to calculate the $^8$B radius and quadrupole moment, the difference in the $^7$Be and $^8$B radii, $r^2({^8{\\rm B}})-r^2({^7{\\rm Be}})$, and the Coulomb displacement energy $E({^8{\\rm Li}})-E({^8{\\rm B}})$ and to study their relation to the $S_{17}(0)$ astrophysical $S$ factor. We find that all these indicators scale linearly with the zero-energy $^7$Be(p,$\\gamma$)$^8$B cross section. Within our three-cluster model we find that the experimentally determined values for the $^8$B radius and quadrupole moment as well as the Coulomb displacement energy is consistently described if the internal cluster size parameters are chosen such that the values for $S_{17}(0)$ are between ($23-25$) eV$\\,$b. This range is more or less compatible with the value currently used in most solar models, $S_{17}(0)=22.4\\pm2.1$ eV$\\,$b \\cite{Johnson}. However, it is slightly inconsistent with the recently adopted new experimental value $S_{17}(0)=19^{+4}_{-2}$ eV$\\,$b \\cite{Adelberger}. As a note of caution we mention that this result has been derived from a linear relation between our four indicators and $S_{17}(0)$ found in our three-cluster model. However, we found that enlarging the $^7$Be model space by adding a $^4{\\rm He}+d+p$ configuration to the $^4{\\rm He}+{^3{\\rm He}}$ configuration increased the $^7$Be quadrupole moment by about $10 \\%$. To investigate the effects which additional configuration might have on the $^8$B properties and in particular on the $S_{17}(0)$ value, requires calculations in model spaces which are beyond $^4{\\rm He}+{^3{\\rm He}}+p$. \\mbox{\\ } The work of A.\\ C.\\ was performed under the auspices of the U.S.\\ Department of Energy. The work has been partly supported by the Danish Research Council and by OTKA grant F019701." + }, + "9802/astro-ph9802121_arXiv.txt": { + "abstract": "Following an approach by Paczy\\'nski \\& Stanek we compare red clump stars with parallaxes known to better than 10\\% in the {\\em Hipparcos}\\/ catalog with the red clump stars observed in three fields in M31 using the {\\em HST}. There are $\\sim 600$ and $\\sim 6,300$ such stars in the two data sets, respectively. The local red clump luminosity function is well represented by a Gaussian with the peak at $M_{I,m}=-0.23$, and the dispersion $\\sigma_{RC}\\approx0.2\\;$mag. This allows a single step determination of the distance modulus to M31 $\\mu_{0,M31} = 24.471\\pm 0.035 \\pm 0.045 \\;$mag (statistical plus systematic error) and the corresponding distance $R_{M31}= 784\\pm 13\\pm 17\\;kpc$. The number of red clump stars is large enough that the formal statistical error in the distance is only $\\lesssim 2$\\%. We also correct the treatment of the local interstellar extinction by Paczy\\'nski \\& Stanek and we obtain the Galactocentric distance modulus $\\mu_{0,GC}=14.57 \\pm0.04 \\pm 0.04\\;$mag (statistical plus systematic error), and the corresponding Galactocentric distance $R_0=8.2 \\pm0.15 \\pm 0.15\\;kpc$. ", + "introduction": "The distance modulus to the M31 galaxy is $\\mu_{0,M31} \\approx 24.4 \\pm0.15\\;$mag (for discussion see e.g.~Huterer, Sasselov \\& Schechter 1995 and Holland 1998). In this paper we follow the approach of Paczy\\'nski \\& Stanek (1998; hereafter: P\\&S) and present an estimate of the distance to M31 based on the comparison between the red clump giants observed locally by the {\\em Hipparcos}\\/ (Perryman et al.~1997) satellite and observed in M31 with the {\\em HST}\\/ (Holland, Fahlman \\& Richer 1996; Rich et al.~1996). These stars are the metal rich equivalent of the better known horizontal branch stars, and theoretical models predict that their absolute luminosity only weakly depends on their age and chemical composition (Seidel, Demarque, \\& Weinberg 1989; Castellani, Chieffi, \\& Straniero 1992; Jimenez, Flynn, \\& Kotoneva 1998). Indeed the absolute magnitude-color diagram of {\\em Hipparcos}\\/ (Perryman et al.~1997, their Figure~3) clearly shows how compact the red clump is. In this paper we determine the variance in the $I$-band magnitude to be only $\\sim 0.15\\;$mag. As discussed by P\\&S, any method of the distance determination which is based on stars suffers from at least four problems: \\begin{enumerate} \\item The accuracy depends on the absolute magnitude determination for the nearby stars; \\item Interstellar extinction has to be determined for the stars in the target source as well as for those near the Sun; \\item The masses, ages, and chemical composition may be different for the stars in the source and for their counterparts near the Sun; \\item The statistical error is large if the number of stars is small. \\end{enumerate} The red clump giants are the only type of stars which do not suffer from the fourth problem. In spite of their large number and sound theoretical understanding these stars have seldom been used as the distance indicators. However, recently Stanek (1995) and Stanek et al.~(1994, 1997) used these stars to map the Galactic bar. P\\&S used the red clump stars observed by OGLE (Udalski et al.~1993) to obtain an estimate of the distance to the Galactic center. In this paper we follow the approach of P\\&S and compare the absolute magnitudes of $\\sim 600$ nearby red clump stars with accurate (better than 10\\%) trigonometric parallaxes measured by {\\em Hipparcos}\\/ with the apparent magnitudes of $\\sim 6,300$ red clump stars observed by the {\\em HST}\\/ in the halo of M31 (Holland et al.~1996) and in the M31 globular cluster G1 (Rich et al.~1996). This comparison gives the distance to M31 in a single step. ", + "conclusions": "" + }, + "9802/astro-ph9802317_arXiv.txt": { + "abstract": "Since the baryon-to-photon ratio $\\eta_{10}$ is in some doubt at present, we ignore the constraints on $\\eta_{10}$ from big bang nucleosynthesis (BBN) and fit the three key cosmological parameters ($h, \\Omega_{\\rm{M}}, \\eta_{10}$) to four other observational constraints: Hubble parameter ($h_{\\rm{o}}$), age of the universe ($t_{\\rm{o}}$), cluster gas (baryon) fraction ($f_{\\rm{o}} \\equiv f_{\\rm{G}}h^{3/2}$), and effective shape parameter ($\\Gamma_{\\rm{o}}$). We consider open and flat CDM models and flat $\\Lambda$CDM models, testing goodness of fit and drawing confidence regions by the $\\Delta\\chi^2$ method. CDM models with $\\Omega_{\\rm{M}} =1$ (SCDM models) are accepted only because we allow a large error on $h_{\\rm{o}}$, permitting $h < 0.5$. Open CDM models are accepted only for $\\Omega_{\\rm{M}} \\gsim 0.4$. $\\Lambda$CDM models give similar results. In all of these models, large $\\eta_{10}$($\\gsim 6$) is favored strongly over small $\\eta_{10}$($\\lsim 2$), supporting reports of low deuterium abundances on some QSO lines of sight, and suggesting that observational determinations of primordial $^4$He may be contaminated by systematic errors. Only if we drop the crucial $\\Gamma_{\\rm{o}}$ constraint are much lower values of $\\Omega_{\\rm{M}}$ and $\\eta_{10}$ permitted. ", + "introduction": "In the context of the hot big bang cosmology, if the number of light-neutrino species has its standard value $N_\\nu =3$, the predicted primordial abundances of four light nuclides (D, $^3$He, $^4$He, and $^7$Li) depend on only one free parameter, $\\eta_{10}$, the universal ratio (at present) of nucleons (baryons) to photons (in units $10^{-10}$). In principle, $\\eta_{10}$ is overdetermined by the observed or inferred primordial abundances of the four light nuclides. Indeed, Steigman, Schramm, and Gunn (1977) have exploited this fact to use big bang nucleosynthesis (BBN) to constrain $N_\\nu$. The status quo ante is that observations, principally of D and $^4$He, have rendered $\\eta_{10}$ one of the best known of the key cosmological parameters: $\\eta_{10}$ = $3.4 \\pm 0.3$ (Walker \\etal, 1991; the error bars being roughly ``1$\\sigma$''). At present, when the microwave background temperature $T$ = 2.728 K (Fixsen \\etal, 1996), the universal baryonic mass-density parameter $\\Omega_{\\rm{B}}\\; (\\,\\equiv 8 \\pi G \\rho_{\\rm{B}}/3H_0^2\\,)$ is related to $\\eta_{10}$ by \\begin{equation} \\Omega_{\\rm{B}}\\,h^2 = 3.667 \\times 10^{-3} \\; \\eta_{10} = 0.0125 \\pm 0.0011. \\label{Eq:Omega_B} \\end{equation} However, recently there have emerged reasons to suspect that $\\eta_{10}$ may not be so well determined, and even that the standard theory of BBN may not provide a very good fit to the current data (Hata \\etal, 1995). There are several options available for resolving this apparent conflict between theory and observation. Although some change in standard physics could offer resolution (e.g., a reduction in the effective value of $N_\\nu$ during BBN below its standard value 3; cf. Hata \\etal, 1995), Hata \\etal\\ (1995) note that large systematic errors may compromise the abundance data (cf. Copi, Schramm, and Turner, 1995). This controversy has been sharpened by new observations giving the deuterium abundances on various lines of sight to high-redshift QSOs. In principle, these data should yield the primordial D abundance, but current results span an order of magnitude. If the low value (D/H by number $\\approx 2 \\times 10^{-5}$; Tytler, Fan, and Burles, 1996; Burles and Tytler, 1996) is correct, then $\\eta_{10} \\approx 7$ in the standard model, but then it seems impossible to reconcile the inferred abundance of $^4$He [Olive and Steigman, 1995 (OS)] with (standard) BBN for this value of $\\eta_{10}$ unless there are large systematic errors in the $^4$He data. If, instead, the high figures (D/H $\\approx 2 \\times 10^{-4}$; Carswell \\etal, 1994; Songaila \\etal, 1994; Rugers and Hogan, 1996) are correct, then D and $^4$He are consistent with $\\eta_{10} \\approx 2$, but modellers of Galactic chemical evolution have a major puzzle: How has the Galaxy reduced D from its high primordial value to its present (local) low value without producing too much $^3$He (Steigman and Tosi, 1995), without using up too much interstellar gas (Edmunds, 1994; Prantzos, 1996), and without overproducing heavy elements (cf. Tosi, 1996, and references therein)? It appears that $\\eta_{10}$, though known to order of magnitude, may now be among the less well-known cosmological parameters. Despite this, large modern simulations which explore other cosmological parameters are often limited to a single value of $\\eta_{10} = 3.4$ (e.g., Borgani \\etal, 1997). Given this unsettled situation Steigman, Hata, and Felten (1997; hereafter SHF) have proposed that it may be constructive to abandon nucleosynthetic constraints on $\\eta_{10}$ entirely and to put $\\eta_{10}$ onto the same footing as the other cosmological free parameters, applying joint constraints on all these parameters based on other (non-BBN) astronomical observations and on theory and simulation. Armed with $\\eta_{10}$ determined in this manner we may then ``predict\" the primordial abundances of the light nuclides and compare with the data to test the consistency of standard BBN. In this contribution to the proceedings of the ISSI Workshop on Primordial Nuclei and Their Galactic Evolution (Bern, Switzerland, 6--10 May 1997) we present a brief description of our approach along with a summary of our results for a ``standard\" (fiducial) choice of the observational constraints. For further details (especially of the many variations on the standard case to be described herein) and references the reader is encouraged to consult SHF. ", + "conclusions": "If BBN constraints on the baryon density are removed (or relaxed), the interaction among the shape-parameter $(\\Gamma)$ constraint, the cluster baryon fraction ($f_{\\rm{G}}$) constraint, and the value of $\\eta_{10}$ assumes critical importance. These constraints still permit a flat CDM model, but only as long as $h < 0.5$ is allowed by observations of $h$. The $f_{\\rm{G}}$ constraint means that large $\\Omega_{\\rm{M}}$ implies fairly large $\\Omega_{\\rm{B}}$. Therefore the exponential term in $\\Gamma$ becomes important allowing $\\Omega_{\\rm{M}} = 1$ to satisfy the $\\Gamma$ constraint. However, values of $\\eta_{10} \\approx 8-15$ are required (see Figures 1 and 2). The best-fit SCDM model has $h \\approx 0.45$ and $\\eta_{10} \\approx 13$, which is grossly inconsistent with the predictions of BBN and the observed abundances of D, $^4$He, and $^7$Li. For $h > 0.5$ a fit to SCDM is no longer possible. The SCDM model is severely challenged. The $\\Gamma$ and age constraints also challenge low-density CDM models. The $\\Gamma$ constraint permits $\\Omega_{\\rm{M}} < 0.4$ only for high $h$, while the age constraint forbids high $h$, so $\\Omega_{\\rm{M}} \\gsim 0.4$ is required. The bound $\\Omega_{\\rm{M}} \\gsim 0.4$ conflicts with the added cluster constraint $\\Omega_{\\rm{o}} = 0.2 \\pm 0.1$ at the 98\\% CL, suggesting strongly that there is additional mass not traced by light. Although a few plausible variations on the CDM models do not affect the constraints very much, removing the $\\Gamma$ constraint would have a dramatic effect. Both high and low values of $\\Omega_{\\rm{M}}$ would then be permitted. The $\\Gamma$ constraint plays a crucial role in our analysis. At either low or high density, the situation remains about the same for the $\\Lambda$CDM models. Because the ages are longer, we can tolerate $\\Omega_{\\rm{M}} \\approx 0.3$ for $h = 0.85$. The $\\Lambda$CDM model therefore accepts more easily the added constraint $\\Omega_{\\rm{o}} = 0.2 \\pm 0.1$. Improved future constraints on $\\Omega_\\Lambda$ will come into play here. Having bounded the baryon density using data independent of constraints from BBN, we may explore the consequences for the light-element abundances. In general, our fits favor large values of $\\eta_{10}$ ($\\gsim 6$) over small values ($\\lsim 2$). While such large values of the baryon density are consistent with estimates from the Ly-$\\alpha$ forest, they do create some tension for BBN. For deuterium there is no problem, since for $\\eta_{10} \\gsim 6$ the BBN-predicted abundance, (D/H)$_{\\rm{P}}$ $\\lsim 3 \\times 10^{-5}$ (2$\\sigma$), is entirely consistent with the low abundance inferred for some of the observed QSO absorbers (Tytler \\etal, 1996; Burles and Tytler, 1996). Similarly, the BBN-predicted lithium abundance, (Li/H)$_{\\rm{P}} \\gsim 2.5 \\times 10^{-10}$ (2$\\sigma$), is consistent with the observed surface lithium abundances in the old, metal-poor stars (allowing, perhaps, some minimal destruction or dilution of the prestellar lithium). However, the real challenge comes from $^4$He where the BBN prediction for $\\eta_{10} \\gsim 6$, Y$_{\\rm{P}} \\gsim 0.248$ (2$\\sigma$), is to be contrasted with the \\hii-region data which suggest Y$_{\\rm{P}}$ $\\lsim 0.238$ (OS, OSS)." + }, + "9802/astro-ph9802067_arXiv.txt": { + "abstract": "The analysis of Perseus data collected with the Medium Energy Concentrator Spectrometer (MECS) on board Beppo-SAX shows that the ratio of the flux of the 8 keV line complex (dominated by Fe K$_{\\beta}$ emission) over the 6.8 keV line complex (dominated by Fe K$_{\\alpha}$ emission) is significantly larger than predicted by standard thermal emission codes. Moreover the analysis of spatially resolved spectra shows that the above ratio decreases with increasing cluster radius. We find that, amongst the various explanations we consider, the most likely requires the plasma to be optically thick for resonant scattering at the energy of the Fe K$_{\\alpha}$ line. We argue that if this is the case, then measures of the iron abundance made using standard thermal emission codes, that assume optically thin emission, can significantly underestimate the true iron abundance. In the case of the core of Perseus we estimate the true abundance to be $\\sim$ 0.9 solar in a circular region with radius of $\\sim 60$ kpc and centered on NGC 1275. Finally we speculate that similar results may hold for the core of other rich clusters. ", + "introduction": "We have accumulated spectra from 5 concentric annuli centered on the emission peak and with bounding radii of 0-2 arcmin, 2-4 arcmin, 4-6 arcmin, 6-8 arcmin and 8-10 arcmin respectively. We have fitted each spectrum in the 3.5-9.5 keV range. We have modeled the spectra with a bremsstrahlung component for the continuum and 2 Gaussian lines, one for the 6.8 keV complex, the other for the 8 keV complex. The ratio of the flux in the lines varies with the radius (see Fig.2), a $\\chi^2$ test rejects a constant line ratio at the 98\\% confidence level. The solid line in Figure 2 is the ratio expected in the case of optically thin thermal emission for temperatures $\\sim$ 5 keV. The dotted line represents the expected ratio when the 8 keV line complex flux in excess of what is expected from an optically thin thermal plasma is distributed as if it came from a point source coincident with the X-ray emission peak. Akimoto et al. (1997) have performed a similar analysis of the line ratio in Perseus using ASCA GIS data. They find a radial gradient which appears to be different from ours. It is rather difficult to make a more detailed comparison, mainly because in their Figure 1: 1) the radii are reported in units of the core radius, but the value of the core radius is not given; 2) errors are not shown so we cannot assess how significant the difference actually is. The same authors present a radial profile of the expected line ratio based on numerical simulations which include the effects of resonant scattering. This profile appears to be in disagreement with our measurements as well as theirs. Even though the authors do not give many details on the parameters chosen for the simulation, from the line ratio predicted in the case of optically thin emission, also shown in their Figure 1, it would seem that they assume a strong temperature gradient within the core of Perseus, possibly to mimic the cooling flow. An analysis of the MECS data shows that, excluding energies below 3 keV where the cooling flow is prominent, the radial temperature gradient in the innermost 200 kpc is rather modest, $\\Delta T \\sim 1.5$ keV. If this is indeed the case then the decrement in the Fe K$_{\\alpha}$ line complex intensity may be considerably larger than that estimated through simulations in Akimoto et al. (1997). We note that explanations of the anomalous line flux ratio involving the ionization equilibrium cannot easily explain the results of the spatially resolved analysis. Indeed if the ionization equilibrium is not thermal because of the active nucleus in NGC 1275, we would expect the ratio to be anomalous only in the innermost spatial bin (the radio lobes extend out to less than 1 arcmin form the nucleus). The data shows that this is not the case. Moreover if the anomalous ratio is the result of an error in the thermal emission codes then, of course, the ratio should be the same throughout the cluster. The interpretation in terms of photon redistribution by resonant scattering appears to be the only viable one among those considered. It is rather difficult to check it quantitatively, as one should consider in detail the abundance and temperature gradients, as well as deviations from a simple King profile for the density due to the cooling flow. The treatment of this problem in real cluster conditions then requires a detailed radiative transfer solution, and is deferred to future works. We here limit ourselves to noting that, according to G86 calculations, the observed factor--of--2 relative depletion of the 6.8 keV complex in the cluster center is consistent with the order--of--magnitude difference in the optical depth of the two complexes. ", + "conclusions": "The analysis of the MECS spectrum of the central region of Perseus has shown that the ratio of the flux in the 8.0 keV line complex over the flux in the 6.8 keV line complex is about a factor 2 larger than expected for an optically thin thermal plasma. Moreover this ratio presents a radial gradient, in the sense that it decreases with increasing cluster radius. We have considered various possible explanations for the behavior of the lines ratio: 1) instrumental effects; 2) a high nickel abundance; 3) a second hot thermal component; 4) non thermal ionization equilibrium; 5) errors in the computation of the collisional ionization equilibrium; 6) resonant scattering of the Fe K$_{\\alpha}$ line. We have argued that resonant scattering is, by far, the most likely explanation for the observed line ratio behavior. We have discussed the implications of an optically thick Fe K$_{\\alpha}$ line. We have found the Fe abundances in the core of Perseus to be significantly larger than previously believed, the true abundances being extremely high ($ \\simgt 0.9 $ solar in a circular region with a radius of $\\sim 60$ kpc centered on the emission peak). We have argued that similar underestimations of the Fe abundance may have also occurred for other rich clusters." + }, + "9802/astro-ph9802251_arXiv.txt": { + "abstract": "We consider massive photon decay reactions via intermediate states of electron--electron-holes and proton--proton-holes into neutrino--antineutrino pairs in the course of neutron star cooling. These reactions may become operative in hot neutron stars in the region of proton pairing where the photon due to the Higgs--Meissner effect acquires an effective mass $m_{\\gamma}$ that is small compared to the corresponding plasma frequency. The contribution of these reactions to neutrino emissivity is calculated; it varies with the temperature and the photon mass as $T^{3/2}m_{\\gamma}^{7/2}{\\displaystyle e}^{-m_{\\gamma}/T}$ for $T< m_{\\gamma}$. Estimates show that these processes appear as extra efficient cooling channels of neutron stars at temperatures $T\\simeq (10^9-10^{10})$~K. ", + "introduction": "The {\\small EINSTEIN}, {\\small EXOSAT} and {\\small ROSAT} observatories measured surface temperatures of certain neutron stars and put upper limits on the surface temperatures of others ( cf. Ref.~\\cite{steuk} and further references therein). Data on the supernova remnants in 3C58, the Crab, and RCW103 indicate rather slow cooling, while the data for Vela, PSR~2334+61, PSR~0656+14, and Geminga point to significantly more rapid cooling. In the so-called standard scenario of neutron star cooling, the most important channel up to temperatures $T\\leq (10^{8} - 10^{9})K$ corresponds to the modified URCA process $n \\, n \\rightarrow n \\, p\\, e\\, \\bar \\nu$. Rough estimates of its emissivity were first made in Ref.~\\cite{BW}. Friman and Maxwell in Ref.~\\cite{fm} recalculated the emissivity of this process in a model, in which the nucleon--nucleon interaction is treated with the help of slightly modified free one-pion exchange. Their result for the emissivity, $\\varepsilon _{\\nu}^{FM}$, proved to be an order of magnitude higher than previously obtained. The value $\\varepsilon _{\\nu}^{FM}$ was used in various computer simulations resulting in the standard cooling scenario; see Ref.~\\cite{T} for example. Subsequent work~\\cite{pion,sv,migrep} took in-medium effects into account in $NN$-interaction, showing that the emissivity of the modified URCA process depends heavily on neutron star mass. For stars of more than one solar mass, the resulting emissivities turned out to be substantially higher than the values given by $\\varepsilon _{\\nu}^{FM}$. These and other in-medium effects were recently incorporated in the computer code in Ref.\\cite{SVSWW} leading to a new scenario of neutron star cooling. For low-mass stars numerical results of the new and standard scenarios more or less coincide. In the present work, we continue to look for enhanced reaction channels. To demonstrate the efficiency of new reaction channels, we compare the results with the emissivity $\\varepsilon _{\\nu}^{FM}$, which dominates cooling in the standard scenario over temperature range under consideration. Besides the modified URCA process, the standard scenario numerical codes also include neutron and proton bremsstrahlung processes $n\\, n\\rightarrow n\\, n \\nu \\bar{\\nu}$ and $n\\, p\\rightarrow n\\, p \\nu \\bar{\\nu}$, which in all models lead to a somewhat smaller contribution to the emissivity than the modified URCA process \\cite{FSB,fm,pion,sv}. Also included are processes that contribute to emissivity in the neutron star crust. These are plasmon decay $\\gamma_{pl}\\rightarrow \\nu\\, \\bar{\\nu}$ \\cite{arw,SB}, electron bremsstrahlung on nuclei $e\\, A\\rightarrow e\\, A\\nu \\bar{\\nu}$ \\cite{P,rf,SB}, electron--positron annihilation $e\\, e^{+}\\rightarrow \\nu \\bar{\\nu}$~\\cite{cm,cs}, and photon absorption by electrons $\\gamma e\\rightarrow e\\, \\nu \\bar{\\nu}$~\\cite{Rit,cs,PBS}. Numerical simulations show that the latter two processes contribute only negligibly to the crust neutrino emissivity at the temperatures under discussion in this paper and they always contribute negligibly to the full neutron star's emissivity; see Fig.~7 of Ref.~\\cite{SB}. When the temperature decreases, it is energetically favorable for neutrons to pair in the neutron star interior and inner crust and for the protons to pair in the star's interior. In a system with nucleon pairing the emissivity of the modified URCA process is suppressed by a factor $\\exp(-(\\Delta_n+\\Delta_p)/T)$ \\cite{fm}, where $\\Delta_n$ and $\\Delta_p$ are the respective neutron and proton gaps, defined by $\\Delta_i(T)=\\Delta_i(0) \\, (T_{c,i}-T) \\, T_{c,i}^{-1} \\, \\Theta(T_{c,i}-T)$ (here $\\Theta(x)$ is the Heaviside step function, $i=\\{p,n\\}$, and $T_{c,i}$ is the corresponding critical temperature for nucleon pairing). At temperatures $T\\ll T_{c,p}, T_{c,n}$ the process becomes marginal. Nevertheless, this star's interior process still dominates those of crust cooling up to temperatures $T\\sim (10^{8} - 10^{9})$~K, depending on the values of the gaps; see Fig.~7 of Ref.~\\cite{SB}. For $T\\leq (1 - 3)\\cdot 10^{8}$~K cooling in the standard scenario is largely dominated by the photon emission from the neutron star surface. In the present work we look for more efficient cooling processes at $T< T_{c,p}, T_{c,n}$. We analyze photon decay into neutrino--antineutrino pairs. The related processes $\\gamma e\\rightarrow e\\, \\nu \\bar{\\nu}$ and $\\gamma p\\rightarrow p\\, \\nu \\bar{\\nu}$ turn out to be suppressed by several orders of magnitude compared to those under discussion, due to the lack of free final states in degenerate fermionic systems, and are therefore not considered here. The contribution of photon decay via electron--electron-hole intermediate states for the case of a normal electron plasma in white dwarfs and neutron star crusts has been calculated by several authors (see Ref.~\\cite{arw} for further references). In an ultrarelativistic electron plasma, a photon acquires an effective in-medium plasmon dispersion law with a gap equal to the electron plasma frequency $\\omega _{pl}\\simeq 2\\,e\\,\\mu_e /\\sqrt{3\\pi}$, where $e$ is the electron charge and $\\mu_e$ denotes the electron chemical potential (we employ units with $\\hbar = c = 1$). Therefore, the contribution to the emissivity of the cited process is suppressed by a factor $\\exp(-\\omega _{pl}/T)$. Nevertheless, in white dwarfs and neutron star crusts, the electron density is not too high, and the process is still effective. In neutron star interiors, the electron density $\\rho_e$ is equal to the proton density $\\rho _p$ by virtue electrical neutrality, and along with $\\beta$ stability one obtains a relation for the total density \\begin{equation} \\label{3} \\rho_e=\\rho _p\\simeq\\, 0.016\\,\\rho_0\\, \\left(\\frac{\\rho }{\\rho_0 }\\right)^2, \\end{equation} where $\\rho_0 \\simeq\\, 0.17$~fm$^{-3}$ denotes the nuclear saturation density, and we use the values of the neutron and proton Fermi momenta~\\cite{fm}, $p_{Fn}\\simeq\\,340(\\rho /\\rho_0)^{1/3}$~MeV and $p_{Fp}=\\mu_e \\simeq\\,85(\\rho /\\rho_0 )^{2/3}$~MeV. Thus, at typical densities for neutron star interiors $\\rho \\stackrel{>}{\\sim} \\rho_0$, the value of the electron plasma frequency is high, e.g., $\\omega _{pl}(\\rho_0)\\approx $ 4.7 MeV for $\\rho \\simeq \\rho_0$, and at temperatures $TT_{c,n}, T_{c,p}$, i.e., in a normal plasma region of the star's crust and interior, photons with approximately the electron plasma frequency\\footnote{A rather small extra contribution also comes from the proton--proton-hole diagram.} $\\omega _{\\rm pl}$ can decay into neutrino pairs, as has been shown in previous estimates~\\cite{arw}. At $T (0.5 - 1)\\rho_0$ the total nucleon--nucleon interaction does not reduce to free one-pion exchange, because of the strong polarization of the medium, whereby a significant part comes from in-medium pionic excitations \\cite{migrep,vs,pion,sv}. Occurring in intermediate states of the reaction, the in-medium pions can also decay into $e\\bar\\nu$, or first into a nucleon--nucleon-hole, which then radiates $e\\bar\\nu$, thereby substantially increasing the resulting emissivity. Other reaction channels such as $n\\rightarrow n_{pair}\\nu\\bar\\nu$ and $p\\rightarrow p_{pair}\\nu\\bar\\nu$ open up in the superfluid phase with paired nucleons~\\cite{flowers,vs,sv}. All these reaction channels give rise to a larger contribution to the emissivity than that of the modified URCA process estimated via free one-pion exchange. Above we compared $\\varepsilon _{\\nu}^{\\gamma}$ with $\\varepsilon _{\\nu}^{FM}$ just because the latter is used in the standard scenarios of neutron star cooling. As we also mentioned in the Introduction, there are other processes like those considered above. Emissivity of the process $p\\gamma_m \\rightarrow p_{pair}\\nu \\bar{\\nu}$ is substantially suppressed (at least by a factor $e^2$ and also due to a much smaller phase-space volume) compared to that of the process $p\\rightarrow p_{pair}\\nu\\bar\\nu$. According to simple estimates, e.g., using Eq. (22) of Ref.~\\cite{Rit}, the process $e\\gamma \\rightarrow e \\nu \\bar{\\nu}$ makes a very small contribution to the emissivity both in the inner crust and in the interior of neutron stars, even when one neglects the photon mass. Thus we may conclude that the process $e\\gamma_m \\rightarrow e \\nu \\bar{\\nu}$ also leads to a minor contribution to the emissivity at the densities and temperatures under consideration. In summary, the processes $\\gamma_m\\rightarrow e\\,e^{-1} + p\\,p^{-1}\\rightarrow \\nu\\bar\\nu$ might be operative over some temperature interval $T\\simeq (10^9-10^{10})$~K, $TF_*)$ is the number of quasars per sky solid angle with flux greater than isotropic flux $F_*$. Then the distribution per cosine of the viewing angle is $\\partial^2 N/\\partial F_*\\partial\\mu = (K/2)F_*^{-\\alpha}$. If we now allow for anisotropy by setting $F_{obs} = F_* f(\\mu)$ (implicitly assuming that the angular radiation pattern is independent of all other quasar characteristics such as luminosity or redshift), the differential source count distribution that we would actually observe is \\begin{equation} {\\partial N \\over \\partial F_{obs}} = \\int_{-1}^{1} \\, d\\mu \\, {\\partial^2 N \\over \\partial F_{obs}\\partial\\mu} = \\int_{-1}^{1} \\, d\\mu \\, {\\partial^2 N \\over \\partial F_{*}\\partial\\mu} {\\partial F_* \\over \\partial F_{obs}} = {K \\over 2} F_{obs}^{-\\alpha}\\int_{-1}^{1} \\, d\\mu \\, f^{\\alpha-1} (\\mu). \\end{equation} For optical fluxes in the neighborhood of $m_B = 16$ -- 17, $\\alpha \\simeq 3.25$ (Hartwick \\& Schade 1990). Since a quasar is only called a BALQSO if it is found within an equatorial wedge defined by $|\\mu| \\leq \\mu_*$, the BAL fraction in a flux-limited sample is \\begin{equation} f_{BAL} = {\\int_{0}^{\\mu_*} \\, d\\mu \\, \\partial^2 N/\\partial F_{obs}\\partial\\mu \\over \\int_{0}^{1} \\, d\\mu \\, \\partial^2 N/\\partial F_{obs}\\partial\\mu} = {\\int_{0}^{\\mu_*} \\, d\\mu \\, f^{\\alpha -1} (\\mu) \\over \\int_{0}^{1} \\, d\\mu \\, f^{\\alpha -1} (\\mu)} . \\end{equation} Thus, if quasars are dimmer in the BAL direction, their population fraction at fixed {\\it observed} flux is diminished because they are submerged in the larger population of intrinsically fainter non-BAL quasars. As already remarked, we cannot compute $f(\\mu)$ with great confidence. However, two plausible guesses will illustrate the possible magnitude of its effect. The first guess is the simplest possible model: that the local limb-darkening is given by the form appropriate to a gray atmosphere in the Eddington approximation (Mihalas 1978), but that the integrated flux is further multiplied by a factor of $\\mu$ because the radiation comes from a flat disk. Then $f_1(\\mu) = |\\mu|(1 + 1.5|\\mu|)$. As a foil to this form, we use a rough analytic fit to detailed radiation transfer calculations which also include a complete treatment of general relativistic corrections (E. Agol, private communication). This analytic fit is $f_2 (|\\mu|) \\simeq 0.23 + 1.54|\\mu|$. The predicted BAL fractions are shown in Fig. 1. Generically, {\\it any} significant limb-darkening in the direction of the absorbing matter leads to a very strong bias against the discovery of BALQSO's. In order to find $\\sim 10\\%$ BALQSO's in an optical flux-limited sample of limb-darkened sources, the true covering fraction $\\mu_* \\simeq 0.5$. When $\\mu_*$ exceeds $\\simeq 0.5$, the fraction of BALQSO's discovered in a flux-limited sample rapidly approaches unity as the limb-darkening weakens and the fraction of non-BALQSO's falls. Conversely, if the absorbing matter is correlated with the {\\it bright} direction, the true fraction might be substantially less than the fraction inferred from flux-limited samples. \\vbox to 4.5in{ \\plotfiddle{\"sampfracs.ps\"}{2.5in}{0}{50}{50}{-20}{-100} \\parbox{4.0in}{ \\vbox to 180pt{\\vfill} \\small\\baselineskip 9pt {\\sc Fig.}~1.---The fraction of BALQSO's that would apear in an optical flux-limited sample for three different limb-darkening laws. The solid curve is the prediction of gray atmosphere limb-darkening on a disk, the dotted curve the approximation to detailed disk models, and the dashed line gives $f_{BAL}$ if the continuum were radiated isotropically. }} This picture also has one more corollary regarding the BAL fraction: if the absorbing matter is in the dim direction, the very faintest quasars (i.e. those at flux levels where $\\alpha$ falls below unity) should be disproportionately BALQSO's. ", + "conclusions": "We have shown that if anisotropy in the continuum emission of quasars is correlated with the direction towards broad absorption line gas, optical flux-limited samples of quasars are severely biassed with respect to the discovery of broad absorption line quasars. If, as the simplest models suggest, quasar continua are limb-darkened at the selection wavelengths, BALQSO's are severely underrepresented in such samples. The inferred covering fraction of broad absorption gas could well be nearer $\\simeq 0.5$ than the usual estimate of $\\simeq 0.1$. Such large covering fractions of absorbing gas would have numerous consequences: Hard X-ray surveys may discover roughly twice as many quasars as would be predicted on the basis of soft X-ray surveys; the ASCA survey may already have done just this. The fact that BALQSO's are somewhat stronger radio sources than the average radio-quiet quasar would also be explained. In addition, resonance scattering in the broad absorption gas could contribute to the emission lines. NV~1240, because it is just 6000~km~s$^{-1}$ redward of Ly$\\alpha$, would be strongly enhanced, providing an alternative explanation [i.e. other than elevated N abundance: Hamann \\& Ferland (1992, 1993)] for the strength of this line in high redshift quasars." + }, + "9802/astro-ph9802119_arXiv.txt": { + "abstract": "During a systematic search for periodic signals in a sample of ROSAT PSPC (0.1--2.4 keV) light curves, we discovered {\\co $\\sim$ 12~min} large amplitude X--ray pulsations in 1WGA~J1958.2+3232, an X--ray source which lies close to the galactic plane. The energy spectrum is well fit by a power law with a photon index of 0.8, corresponding to an X--ray flux level of $\\sim 10^{-12}$~erg~cm$^{-2}$~s$^{-1}$. The source is probably a long period, low luminosity X--ray pulsar, similar to X Per, or an intermediate polar . ", + "introduction": "Recent ROSAT observations have substantially increased the number of known X--ray pulsators (XRPs; Hughes 1994; Dennerl, Haberl \\& Pietsch 1995; McGrath \\etal 1994; Israel \\etal 1997a, 1997b) and cataclysmic variables (CVs; Reinsch \\etal 1994; Shafter \\etal 1995; Buckley \\etal 1995; Haberl \\& Motch 1995; Burwitz \\etal 1996; Friedrich \\etal 1996; Singh \\etal 1996). In this paper we report the discovery of highly significant pulsations at a period of {\\co 12\\,min} in the X--ray flux of 1WGA\\,J1958.2+3232, a serendipitous ROSAT PSPC source which lies in the galactic plane. We also discuss the possible nature of the compact object responsible for these pulsations. ", + "conclusions": "We discovered large amplitude {\\co 12\\,min} pulsations in the X--ray flux of 1WGA~J1958.2+3232. This periodicity likely originates from polar cap accretion onto a rotating magnetic compact star, either a white dwarf in an intermediate polar, or a neutron star in an X--ray binary system. In the latter case 1WGA~J1958.2+3232 would be one of the very few XRPs known with a spin period of $> 500$ s and a low X--ray luminosity. If instead 1WGA~J1958.2+3232 hosts an accreting magnetic white dwarf rotation, it would likely represent a distant member of the class of low galactic latitude intermediate polars, recently discovered in the ROSAT all--sky survey. The additional possibility that the X--ray modulation arises from the orbital motion of a high inclination low mass X--ray binary with an accretion disk corona cannot be excluded at present." + }, + "9802/astro-ph9802269_arXiv.txt": { + "abstract": "We investigate the large-scale clustering of radio sources in the FIRST 1.4-GHz survey by analysing the distribution function ({\\it counts in cells}). We select a reliable sample from the the FIRST catalogue, paying particular attention to the problem of how to define single radio sources from the multiple components listed. We also consider the incompleteness of the catalogue. We estimate the angular two-point correlation function $w(\\theta)$, the variance $\\Psi_2$, and skewness $\\Psi_3$ of the distribution for the various sub-samples chosen on different criteria. Both $w(\\theta)$ and $\\Psi_2$ show power-law behaviour with an amplitude corresponding a spatial correlation length of $r_0 \\sim 10 h^{-1}$Mpc. We detect significant skewness in the distribution, the first such detection in radio surveys. This skewness is found to be related to the variance through $\\Psi_3=S_3(\\Psi_2)^{\\alpha}$, with $\\alpha=1.9\\pm 0.1$, consistent with the non-linear gravitational growth of perturbations from primordial Gaussian initial conditions. We show that the amplitude of variance and skewness are consistent with realistic models of galaxy clustering. ", + "introduction": "Surveys of optical and infrared galaxies have revealed the rich structure of the universe and have provided much information on the large-scale structure out to redshifts $z\\sim 0.1$ (e.g. APM, Maddox et al., 1990 and IRAS, Fisher et al., 1993). In contrast radio sources, although representing only a small fraction of all galaxies, can be detected over significant cosmological distances (up to $z\\sim 4$), sampling much larger volumes of space and therefore with the potential to provide information on much larger physical scales. It had been suggested (Kaiser, 1984) that galaxies form preferentially in high-density peaks of the underlying mass distribution. If it is true, then the statistics of galaxy distributions provide us with information, although biased, about the underlying matter distribution. It has long been debated whether radio galaxies are clustered or isotropic on the largest scale. The study by Webster (1976), which looked at 8000 radio sources, found $<3\\%$ variability in the number of sources in randomly-placed 1-Gpc cubes. This led to a widely-accepted view that radio sources were isotropically distributed. Even if this were not true, the large range in intrinsic luminosity of radio sources might effectively wash out structures when the distribution is projected onto the sky with information about radial distribution effectively lost. Other studies (Seldner \\& Peebels (1981), Shaver \\& Pierre (1989)) reported a detection of slight clustering of nearby radio sources, while Benn \\& Wall (1995) and Baleisis et al. (1998) discussed other measures of anisotropy in radio surveys. Clustering in the 4.85 Ghz Green Bank (87GB: Gregory \\& Condon, 1991) and in the Parkes-MIT-NRAO (PMN: Griffith \\& Wright, 1993) catalogues was studied by correlation-function analysis (Kooiman et al., 1995; Sicotte, 1995; Loan, Wall \\& Lahav, 1997). These studies indicated that radio objects are actually more strongly clustered than local optically-selected galaxies. This conclusion was confirmed by correlation analysis of the FIRST survey (Cress et al., 1996). One of the possible ways of investigating clustering properties of radio sources is by means of the distribution function ({\\it counts in cells}) i.e. the probability of finding $N$ galaxies in a cell of particular size and shape. This analysis includes all the moments of the distribution function and therefore provides a more complete description of large-scale structure. Furthermore it can be shown (Peebles, 1980; Frieman \\& Gazta\\~naga, 1994) that the higher-order moments of the galaxy distribution can be used as a test of non-linear models for large-scale structure. This paper presents a counts-in-cells analysis carried out for the FIRST radio survey. We focus on the second and third moment of the distribution together with the angular two-point correlation function of the sample, and we test the predictions of different cosmological models. We report a detection of skewness $S_3 =\\Psi_3/\\Psi_2^2=const$ (with $\\Psi_2$ and $\\Psi_3$ defined as the second and the third moment of the distribution). Our measurement accords with the hypothesis of non-linear growth of observed structures by gravitational clustering from initially-Gaussian density fluctuations. In Section 2 we describe the catalogue in its original form and explain the procedures providing us with modified samples analysed in the rest of the paper. In Section 3 we present the results of our analysis for the angular two-point correlation function, and the angular second and third moments of the distribution. Section 4 discusses the deprojection of our 2-d measurements to estimate the quantities describing spatial distribution. Section 5 summarises our conclusions. ", + "conclusions": "" + }, + "9802/astro-ph9802096_arXiv.txt": { + "abstract": "We calculate the relativistic thermal bremsstrahlung Gaunt factor for the high-temperature plasma which exists in clusters of galaxies. We calculate the Gaunt factor by employing the Bethe-Heitler cross section corrected by the Elwert factor. We also calculate the Gaunt factor by using the Coulomb-distorted wave functions for nonrelativistic electrons following the method of Karzas and Latter. By comparing the Gaunt factors calculated by these two different methods, we carefully assess the accuracy of the calculation. We present the numerical results in the form of tables. ", + "introduction": "High-temperature plasmas exist in the clusters of galaxies (Arnaud et al. 1994; Markevitch et al. 1994; Markevitch et al. 1996; Holzapfel et al. 1997). Some clusters have extremely high-temperature electrons, $k_{B} T_{e}$ = $10 \\sim 15$keV. Relativistic expressions for the thermal bremsstrahlung emissivity have been discussed by many authors (Gould 1980; Rephaeli \\& Yankovitch 1997). However, the relativistic expressions have been so far derived by power-series expansions. In view of the recent advance in the accuracy of the observation of the electron temperatures in the clusters of galaxies, it appears worthwhile to assess the accuracy of the relativistic expressions for the thermal bremsstrahlung emissivity. Concerning the inverse thermal bremsstrahlung (free-free absorption) Itoh and his collaborators have calculated the relativistic Gaunt factor (Itoh, Nakagawa, \\& Kohyama 1985; Nakagawa, Kohyama, \\& Itoh 1987; Itoh, Kojo, \\& Nakagawa 1990; Itoh et al. 1991; Itoh et al. 1997). The main purpose of these papers was to present accurate opacities for dense high-temperature stellar plasmas. Therefore, high-density regimes were treated in these papers. In the present paper we will deal with the high-temperature, low-density regime which is relevant to the hot gas in the clusters of galaxies. In calculating the relativistic thermal bremsstrahlung Gaunt factor for this regime, we will make use of the Bethe-Heitler cross section corrected by the Elwert factor. In order to assess the accuracy of the present calculation, we will also calculate the Gaunt factor using the Coulomb-distorted wave functions for nonrelativistic electrons following the method of Karzas and Latter (1961). The present paper is organized as follows. We will give formulations for the calculation of the relativistic thermal bremsstrahlung Gaunt factor in $\\S$ 2. The numerical results will be presented in $\\S$ 3. We will discuss the results and give concluding remarks in $\\S$ 4. ", + "conclusions": "" + }, + "9802/astro-ph9802213_arXiv.txt": { + "abstract": "The $z = 2.286$ IRAS galaxy F10214+4724 remains one of the most luminous galaxies in the Universe, despite its gravitational lens magnification. We present optical and near-infrared spectra of F10214+4724, with clear evidence for three distinct components: lines of width $\\sim 1000 ~ \\rm km s^{-1}$ from a Seyfert-II nucleus; $\\stackrel{<}{_\\sim}200$ km s$^{-1}$ lines which are likely to be associated with star formation; and a broad (\\mbox{$\\sim4000\\rm~km~s^{-1}$}) C{\\sc iii}]~$1909$\\AA\\ emission line which is blue-shifted by \\mbox{$\\sim1000\\rm~km~s^{-1}$} with respect to the Seyfert-II lines. Our study of the Seyfert-II component leads to several new results, including: (i) From the double-peaked structure in the Ly$\\alpha$ line, and the lack of Ly$\\beta$, we argue that the Ly$\\alpha$ photons have emerged through a neutral column of $N_{\\rm H} \\sim 2.5 \\times 10^{25} \\rm ~m^{-2}$, possibly located within the AGN narrow-line region as argued in several high redshift radiogalaxies. (ii) The resonant O{\\sc vi} $1032$, $1036$\\AA\\ doublet (previously identified as Ly$\\beta$) is in an optically thick (\\mbox{1:1}) ratio. At face value this implies an an extreme density ($n_{e} \\sim 10^{17} ~ \\rm m^{-3}$) more typical of broad line region clouds. However, we attribute this instead to the damping wings of Ly$\\beta$ from the resonant absorption. (iii) A tentative detection of HeII $1086$ suggests little extinction in the rest-frame ultraviolet. ", + "introduction": "\\label{sec:introduction} The \\mbox{$z=2.286$} IRAS galaxy FSC 10214+4724 is one of the most apparently luminous objects in the Universe, and its discovery (Rowan-Robinson et al. 1991) led to much speculation about its possible status as a protogalaxy. This speculation was based on the extreme bolometric luminosity of the object (Rowan-Robinson et al. 1991), and more specifically on the huge gas mass and star formation rate inferred from the sub-mm molecular line and continuum detections (e.g. Solomon, Downes \\& Radford 1992; Rowan-Robinson et al. 1993). This speculation was dampened by a series of papers which proved that F10214+4724 is being gravitationally lensed, and, at all wavebands, is intrinsically an order of magnitude or more dimmer than it first appeared. Although a lensing bias was suspected by a number of authors ({\\it e.g.} Elston {\\it et al.} 1994, Trentham 1995), the first direct empirical evidence of strong gravitational lensing was provided by a deep near-infrared image (Matthews et al. 1994) which revealed an arc-like structure centred on a galaxy close to the line of sight to F10214+4724. Several sets of authors published lensing interpretations (Graham \\& Liu 1995; Serjeant {\\it et al.} 1995; Broadhurst \\& Leh\\'ar 1995) which were confirmed by the appearance of the HST image of Eisenhardt {\\it et al.} (1996): this image contained highly elliptical, high surface brightness features characteristic of strong lensing, as well as a clear counter-image. These papers also attempted to constrain the redshift of the system responsible for the gravitational lensing, our contribution (Serjeant {\\it et al.} 1995) being spectroscopy of two galaxies projected $\\approx 1$ and $\\approx$3 arcsec from F10214+4724. This work revealed tentative $4000$\\AA\\ breaks at \\mbox{$z\\simeq0.90$} in both galaxies, later confirmed using fundamental plane arguments in HST imaging (Eisenhardt {\\it et al.} 1996), and also tentatively supported by a weak absorption line at $z=0.893$ in the F 10214+4724 spectroscopy of Goodrich et al. (1996). The HST R-band image of F10214+4724 implies magnifications of $\\sim100$ (Eisenhardt {\\it et al.} 1996) in this waveband, but it now seems likely that differential flux magnification causes lower magnification factors for the more extended structure, as argued by several authors. Gravitational lensing appeared to offer a compelling explanation for the extreme luminosity of FSC 10214+4724 ({\\it e.g.}, Broadhurst \\& L\\'{e}har 1995). As a result, the IRAS galaxy is no longer so extreme in its properties: indeed, in many respects, it resembles local ultraluminous infrared galaxies and Seyfert-II galaxies. Nevertheless, both Downes {\\it et al.} (1995) and Green \\& Rowan-Robinson (1996) argue for a bolometric magnification factor of $\\stackrel{<}{_\\sim}10$ using arguments based on minimum black body sizes. FSC 10214+4724 remains one of the most intrinsically luminous objects in the Universe. The importance in this object still lies in studying whether high-$z$ hyperluminous activity, such as that seen in FSC 10214+4724, differs in all but scale from local objects, and in determining the relative contributions of the starburst and AGN components (Elston et al. 1994; Lawrence et al. 1993, 1994; Soifer et al. 1995; Goodrich et al. 1996; Kroker et al. 1996; Hughes, Dunlop \\& Rawlings 1997). Prior to making the observations reported in this paper there had been no direct evidence for the presence of an embedded broad-line (e.g. Seyfert-I or quasar) nucleus in $\\rm F10214+4724$, although high ($\\approx$20 per cent) rest-frame ultraviolet polarization (Lawrence et al. 1993; Jannuzi et al. 1995) suggested that one was present. This situation changed with the deep spectropolarimery of Goodrich et al. (1996) showing clear broad lines in polarized light. This extended the close spectral similarities between $\\rm F10214+4724$ and Seyfert-II galaxies, specifically NGC1068, which was first remarked on by Elston et al. (1994). Prior to our observations there had also been no reported detection of the narrow ($\\stackrel{<}{_\\sim}200$ km s$^{-1}$) emission lines expected from any star-forming activity in $\\rm F10214+4724$. A star-forming component is expected if the analogy with NGC1068 is to be complete. This situation also changed during the preparation of this paper. Using a novel imaging near-infrared spectrometer Kroker et al. (1996) presented evidence for spatially-extended narrow H$\\alpha$ emission just as expected if the Seyfert-II nucleus of F10214+4724 is accompanied by a circumnuclear starburst. In this paper we present, analyse and interpret optical and near-infrared spectroscopy of IRAS FSC 10214+4724. The details of data acquisition and analysis are given in Section 2. In Section 3 we compare our results with previous and contemperaneous spectroscopic studies of F10214+4724. In Section 4 we interpret the data on the Seyfert-II emission line region including a discussion of optical depth effects on the resonance lines, and some modelling of the spectrum using the photoionisation code {\\sc CLOUDY} ({\\it e.g.} Ferland 1993, 1996). In this section we reach conclusions about the Seyfert-II properties of F10214+4724 which differ significantly from those reached by previous studies. In Section 5 we interpret the data on the region responsible for the narrow ($\\stackrel{<}{_\\sim}200 ~ \\rm km s^{-1}$) H$\\alpha$ line seen in our near-infrared spectrum: this feature is likely to be a signature of star formation. In Section 6 we pass some concluding remarks on the nature of $\\rm F10214+4724$. Further data on galaxies foreground to $\\rm F10214+4724$ are given in Appendix A, and our attempts at identifying the weak emission line at $2067$\\AA\\ are discussed in Appendix B. ", + "conclusions": "The presence of the optically thick O{\\sc vi} $1032,1037$\\AA\\ doublet appears at face value to imply extremely high densities for narrow emission line gas, $n_{\\rm H}\\stackrel{>}{_\\sim}10^{17}$ m$^{-3}$. However, we argue that it is more easily attributable to the Ly$\\beta$ damping wings of resonant scattering material, probably withing the AGN narrow line region. Differential magnification appears to play a significant role in the emission line spectrum of F10214+4724, in which we identify three distinct kinematic components: quasar broad lines ($\\sim 4000$ km s$^{-1}$), Seyfert II narrow lines ($\\sim 1000$ km s$^{-1}$) and a starburst component ($\\stackrel{<}{_\\sim}200$ km s$^{-1}$). The density, ionisation parameter, number and total mass of Seyfert II narrow line clouds all resemble local Seyferts, and our interpretation of the resonant scattering agrees with that of Villar-Martin {\\it et al.} (1996) for high-redshift radiogalaxies. The flux from the narrowest H$\\alpha$ component is in excellent agreement with radiative transfer models which comprise similar quasar and starburst bolometric contributions." + }, + "9802/astro-ph9802025_arXiv.txt": { + "abstract": "We propose a simple and unambiguous way to deduce the parameters of black holes which may reside in AGNs and some types of X-ray binaries. The black-hole mass and angular momentum are determined in physical units. The method is applicable to the sources with periodic components of variability, provided one can assume the following: (i)~Variability is due to a star or a stellar-mass compact object orbiting the central black hole and passing periodically through an equatorial accretion disk (variability time-scale is given by the orbital period). (ii)~The star orbits almost freely, deviation of its trajectory due to passages through the disk being very weak (secular); the effect of the star on the disk, on the other hand, is strong enough to yield observable photometric and spectroscopic features. (iii)~The gravitational field within the nucleus is that of the (Kerr) black hole, the star and the disk contribute negligibly. ", + "introduction": "There is an ever-growing theoretical and observational evidence that the dark masses present in some galactic nuclei and X-ray binaries are black holes (Rees 1998). Several ways have been suggested of how to deduce the properties of these black holes. The first rough estimates of their mass were based on energy considerations and limits implied by the shortest variability timescale (e.g.\\ Begelman, Blandford, \\& Rees 1984). More precise methods became possible with modern observational techniques like HST, VLBI and X-ray satellites. The aim of the present paper is to propose and further discuss a method to determine the parameters of a black hole in a system with two periodic components in the observed signal which are due to orbital motion of a stellar object and its passages through the accretion disk of the hole. We start with a brief summary of the subject. In galactic nuclei, the existence and size of the dark central mass are deduced from the curve of the central surface brightness of the nuclei and from the spatial distribution and dynamics of surrounding gas and stars (see Kormendy \\& Richstone 1995 for a review and references). In particular, the standard model for active galactic nuclei (AGNs) with a black hole and an accretion disk (e.g.\\ Rees 1984; Blandford \\& Rees 1992) has been further supported when the nuclear gas in several active galaxies was found to be in a gravity-dominated nearly circular motion in a disk (e.g.\\ Jaffe et al.\\ 1993; Gallimore, Baum, \\& O'Dea 1997). Rotation curves of the observed nuclear gas disks indicate a central compact dark mass of some (2.0--3.5)$\\times$$10^{9}M_{\\odot}$ in M\\,87 (NGC 4486) (Merritt \\& Oh 1997; Macchetto et al.\\ 1997), of about 4.9$\\times$$10^{8}M_{\\odot}$ in NGC 4261 (Ferrarese, Ford, \\& Jaffe 1996), and of 3.6$\\times$$10^{7}M_{\\odot}$ in NGC 4258 (Miyoshi et al.\\ 1995; Maoz 1995). The observed excess of quasars at high redshifts has stimulated a search for black holes in quiescent galaxies, since they may have been both the engine and then the residue of a former activity (Haehnelt \\& Rees 1993). In this case the most reliable present data come from the observations of stellar motion (van den Bosch \\& de Zeeuw 1996). Velocity profiles of several stellar nuclear disks offer strong evidence for black holes --- of 2.6$\\times$$10^{6}M_{\\odot}$ in our own Galaxy (Genzel et al.\\ 1997; cf.\\ Mezger, Duschl, \\& Zylka 1996), of 2$\\times$$10^{9}M_{\\odot}$ in NGC 3115 (Kormendy et al.\\ 1996a), of 1$\\times$$10^{9}M_{\\odot}$ in NGC 4594 (the Sombrero galaxy) (Kormendy et al.\\ 1996b), of 3$\\times$$10^{6}M_{\\odot}$ in M\\,32 (van den Marel et al.\\ 1997), and probably also of 6$\\times$$10^{8}M_{\\odot}$ in NGC 4486B (Kormendy et al.\\ 1997). However, current optical and radio studies still probe only regions above $10^{4}\\div 10^{5}$ gravitational radii of the putative holes which is by far not enough to resolve any imprints of the relativistic effects due to the collapsed centre, and in particular to deduce the rate of its rotation. Information from within the region of a few tens of gravitational radii have been brought by X-ray satellite observations of a sample of Seyfert 1 galaxies (Nandra et al.\\ 1997 and references therein; see also Fabian 1997 for a survey, and Tanaka et al.\\ 1995 for the best-quality measurement from MCG-6-30-15). These observations yielded the profiles of the fluorescent iron K$\\alpha$ emission line which is likely to be induced by illumination of the very inner parts of an accretion disk by a source off the disk plane (e.g.\\ Matt, Perola, Piro, \\& Stella 1992; Petrucci \\& Henri 1997). Indeed, the observed broad profile skewed to lower energies is best explained in terms of the Doppler and gravitational redshifts (Fabian et al.\\ 1995); the exact origin of the illuminating source remains unclear. Evidence for smallness of the emitting regions also comes from the resolved rapid variability of the K$\\alpha$ lines (Nandra et al.\\ 1997, and references therein). Note that recently Bao et al.\\ (1997) reminded another unique (and measurable) feature of black hole sources involving accretion disk: the energy-dependent variability of polarization of their X-rays, originally discussed by Stark \\& Connors (1977). A way to estimate parameters of the central black hole is strongly related to the source variability. It is assumed, in fact, that a time delay between variations of the emission line strength of the AGN and that of the continuum can be due to the time travel between the central source and the surrounding line-emitting gas. Radial distance and velocity of propagation of the disturbances impose limits on the black-hole mass (Blandford \\& McKee 1982; Krolik et al.\\ 1991), but the resulting estimates inherit uncertainties of the underlying assumptions. Alternatively, it has been proposed to deduce the masses of the central nuclear bodies from temporal changes of the observed emission lines (Stella 1990), while rotational parameters from the position, intensity, width (Hameury, Marck, \\& Pelat 1994; Martocchia \\& Matt 1996), and profile of the lines (Bromley, Chen, \\& Miller 1997; Dabrowski et al.\\ 1997; Reynolds \\& Begelman 1997; Rybicky \\& Bromley 1998; cf.\\ also Bao, Hadrava, \\& {\\O}stgaard 1994 and 1996). The X-ray photometric light-curve profile from a hot spot orbiting not far from the horizon also reflects rotation of the black hole (Asaoka 1989; Karas 1996). Similar research has been focused on galactic X-ray binaries where a central engine was also proposed, namely an accretion disk fed by overflow from a secondary star. Here the evidence for a black hole involves non-stellar appearance (``invisibility'') of one of the components, the presence of X-radiation, variability, spectral features (mainly the presence of relatively strong ultrasoft component and of hard X-ray tail), and of course a large lower limit on mass of the dark component ($\\gtrsim3M_{\\odot}$), deduced from the orbital parameters of the (stellar) companion (namely from the mass function of the system). (See the surveys in Lewin, van Paradijs, \\& van den Heuvel 1995; the reviews of black-hole binaries with a thorough list of references have also been given by Barret, McClintock, \\& Grindlay 1996; Charles 1997. Later results are due to Beekman et al.\\ 1996 and 1997; Orosz \\& Bailyn 1997.) It has been proposed recently to infer the parameters of the dark component from observable behaviour of the disk. An observable modulation of the X-ray emission may be caused by disk vibrations whose lowest frequency depends on the mass and rotation of the hole (Nowak et al.\\ 1997 and references therein). Zhang et al.\\ (1997) proposed, also on the basis of the standard thin accretion disk model, how the angular momentum (spin) of the black hole in an X-ray binary could be inferred from the strength of an ultrasoft X-component. Very recently Cui, Zhang, \\& Chen (1998) have proposed ``that certain types of quasi-periodic oscillations (QPO) observed in the light-curves of black-hole binaries are produced by X-ray modulation at the precession frequency of accretion disks, because of relativistic dragging of inertial frames around spinning black holes''. Given the mass of the hole, they are able to derive its spin by comparing the computed disk precession frequency with that of the observed quasi-periodic oscillations. This mechanism requires a warped accretion disk, the assumption that clearly calls for further investigation (Markovi\\'c \\& Lamb 1998). Scheme for the source variability discussed in the present paper is different from that which has been proposed for QPOs in the above-mentioned papers. Here, an orbiting companion of the black hole is involved --- the assumption which restricts variability time-scales relevant for the model. Narayan, McClintock, \\& Yi (1996), and Narayan, Garcia, \\& McClintock (1997) argued that the black-hole X-ray binaries could also be recognized according to a larger variation in luminosity between their bright and faint states than is expected in the sources with neutron stars. This way of identifying black holes follows from the low radiative efficiency of ad\\-vec\\-t\\-ion-do\\-mi\\-na\\-ted accretion flows which could occur around black holes (with greater ``sucking-in'' power and no solid surface) rather than around other compact objects. Various under-luminous accretion-powered astrophysical systems, mainly quiescent transient X-ray sources and low-luminosity galactic nuclei, could be interpreted as black holes with advection-dominated accretion disks (see Lasota \\& Abramowicz 1997 for a survey); Reynolds et al.\\ (1996) suggested that advection-dominated mode of the final stages of accretion could also account for the ``quiescence'' of the black hole in M\\,87. The lack of a hard surface of a black hole also plays crucial role in an alternative explanation by King, Kolb, \\& Szuszkiewicz (1997) of X-ray transients which are dynamical black-hole candidates, namely in terms of a weaker (stabilizing) irradiation of the disk by the central accreting source. In the present paper we propose a method applicable, under certain assumptions, to the black-hole sources with periodic components of variability. It provides, unambiguously, {\\em{}both the mass and the specific angular momentum of the black hole in physical units}. We start from the model of possible periodic variability of black-hole sources which considers a thin accretion disk lying in the hole's equatorial plane, and an object (a ``star'') which intersects the disk periodically while orbiting about the centre.\\footnote {See e.g.\\ Krolik et al.\\ 1991; Mineshige, Ouchi \\& Nishimori 1994; Ipser 1994; Zakharov (1994); Bao, Hadrava, \\& {\\O}stgaard (1994); Kanetake, Takeuti, \\& Fukue (1995), and references therein for alternative explanations of periodic or quasi-periodic variability.} Such a system is characterized by two angular frequencies --- that of azimuthal revolution of the star and that of its latitudinal oscillation about the equatorial plane. Both frequencies are in principle measurable at infinity, the azimuthal one from spectrophotometry, while the latitudinal one from photometry provided that the passages of the star produce a strong enough modulation of the source. Karas \\& Vokrouhlick\\'y (1994) illustrated, by Fourier analysis of simulated photometric data, how the respective two peaks can be recognized in the power spectrum. This model was cultivated notably in connection with the NGC~6814 galaxy whose putative variability, however, was later recognized as being due to a source in our Galaxy. Other possible targets are discussed in the current literature. For example, optical outbursts in the blazar OJ~287 have recently been modelled in terms of a black-hole binary system by Villata et al.\\ (1998). (The source exhibits several time-scales: feature-less short-term variability, 12-yr cycle, and, possibly, a 60-yr cycle.) These authors, however, presume both components to be of comparable masses while in our calculation frequencies are determined under assumption that the secondary is much less massive than the primary (cf. also Sundelius et al.\\ 1997). As another example, a 16-hr periodicity in the X-ray signal from the Seyfert galaxy IRAS 18325-5926 was described by Iwasawa et al.\\ (1998). Let us suppose that the disk has only a weak dynamical influence on the star and that stellar tides are negligible (the star is assumed much smaller than the typical curvature radius of the field around) as well as gravitational radiation and self-gravity both of the disk and of the star itself. Then the worldline of the star is very close to a geodesic in a pure gravitational field of the central rotating black hole. This is described by the Kerr metric which reads (Misner, Thorne \\& Wheeler 1973, p.\\ 878), in Boyer-Lindquist spheroidal coordinates ($t,r,\\theta,\\phi$), in geometrized units (in which $c=G=1$, $c$ being the speed of light in vacuum and $G$ the gravitational constant) and with the ($-$+++) signature of the metric tensor ($g_{\\mu\\nu}$), \\begin{eqnarray} \\label{metric} {\\rm d}s^{2} & = & -\\frac{\\Delta\\Sigma}{\\cal A}\\,{\\rm d}t^{2} +\\frac{{\\cal A}}{\\Sigma}\\,\\sin^{2}\\theta\\; ({\\rm d}\\phi-\\omega_{\\rm K}{\\rm d}t)^{2} \\nonumber \\\\ & & +\\frac{\\Sigma}{\\Delta}\\,{\\rm d}r^{2} +\\Sigma{\\rm d}\\theta^{2}, \\end{eqnarray} where $M$ and $a$ denote mass and specific rotational angular momentum of the source and \\begin{equation} \\Delta=r^{2}-2Mr+a^{2}, \\;\\;\\; \\Sigma=r^{2}+a^{2}\\cos^{2}\\theta, \\end{equation} \\begin{equation} {\\cal A}=(r^{2}+a^{2})^{2}-\\Delta a^{2}\\sin^{2}\\theta, \\;\\;\\; \\omega_{\\rm K}=2Mar/{\\cal A}. \\end{equation} It has been demonstrated (Syer, Clarke, \\& Rees 1991; Vokrouhlick\\'y \\& Karas 1993, 1998) that under the above-described circumstances the stellar orbit undergoes three secular changes: a decrease of the semi-major axis (the star spirals towards the centre due to the loss of energy in collisions with the disk), circularization (the orbit becomes spherical, $r={\\rm const}$), and a decrease of the amplitude of precession of the orbital plane about the equatorial plane of the centre (the orbit gradually declines into the disk plane). Since the time scale for circularization is typically found shorter than, or of the same order as, the time scale necessary to drag the orbit into the disk, one can expect that at late stages of evolution of the hole-disk-star system the star follows a nearly equatorial spherical geodesic. This stage is also the one in which the star stays for a relatively long time. In the next section we discuss relevant properties of a precessing orbit in the Kerr spacetime. Then, in Sect. \\ref{eqset}, we derive simple relations which appear to be appropriate for practical study of the relevant objects. Equations further simplify if the object orbits not very close to the black hole, as discussed in Sect. \\ref{simple}. Our method employs also the observed emission-line profiles which were vastly studied in recent literature; we summarize the relevant formulae and results in Appendix~A. ", + "conclusions": "In the present paper, we outlined general features of our method; an application to particular sources will be discussed separately (work in progress). We will conclude by brief comments on several important points. First, although the method described above {\\em{}in principle\\/} yields {\\em{}all\\/} the relevant quantities, it might be combined advantageously with independent determination of some of the parameters. In particular, knowing $M$, one can deduce $a/M$ and $r/M$ from eqs.\\ (\\ref{a}) and (\\ref{r}). Observed emission-line profiles from relativistic accretion disks around black holes have recently attracted great interest and they play a significant role also in our method. Apart from the well-known double-horn feature (Laor 1991), characteristics of the lines depend on rather uncertain properties of accretion flows, e.g.\\ on advection velocity (Fukue \\& Ohna 1997), limb-darkening law (Rybicki \\& Bromley 1997) and shape of the disk (Pariev \\& Bromley 1997). One therefore needs to investigate a broad range of models to determine widths, required in our model, and to reject unacceptable profiles. We present some results in this direction in the Appendix. Crucial point in our considerations is the (quasi)-periodic variability of the black-hole--disk source. Periodic modulation of the accretion flow is the most likely signature for a stellar companion close to the central black hole (Podsiadlowski \\& Rees 1994). Each star's passage through the accretion disk pulls some amount of gaseous material out of the disk (Zurek, Siemiginowska, \\& Colgate 1994). This material temporarily covers the innermost region of the disk and modulates the observed radiation (both continuum and line).\\footnote {Syer \\& Clarke (1995) discussed a response of the disk to a body moving completely inside and determined conditions under which a gap can develop and survive (see also \\v{S}lechta 1998).} It can thus be anticipated that a variable signal is recognized in the X-ray band, however, details of the mechanism remain rather uncertain. Note that also the star may be affected by the collisions and tidal forces considerably --- it may lose its outer layers or even become disrupted (Frank \\& Rees 1976; Marck, Lioure, \\& Bonazzola 1996; Diener et al.\\ 1997; Loeb \\& Ulmer 1997). A stripped stellar core is an intense source of (ir)radiation which may survive for many orbital periods (Rees 1998). In any case, we assumed here that the star is only {\\em weakly\\/} affected by the disk, i.e.\\ that its orbit may well be approximated by a nearly equatorial spherical geodesic. Radiation from the presumed accretion disks can be modulated by their precession, vibrations and various instabilities. It is also very well possible that the disks are covered by numerous irregularities (phenomenologically designated as bright spots), contributing to the featureless short-term X-ray variability which is indeed observed (Abramowicz et al.\\ 1991; Mangalam \\& Wiita 1993). Then our scheme must anticipate that well-separated, profound irregularities develop on the disk surface (due to stellar passage) and survive several orbital periods. Although no AGN or black-hole X-ray binary with clear periodic variability is currently known, the cases of rapidly variable K$\\alpha$ line (Iwasawa et al.\\ 1996; Yaqoob et al.\\ 1996), and that of quasi-periodic AGNs (Papadakis \\& Lawrence 1993;\\footnote {However, cf.\\ Tagliaferri et al.\\ (1996).} Lehto \\& Valtonen 1996; Stothers \\& Sillanp\\\"a\\\"a 1997; Valtonen \\& Lehto 1997) and X-ray sources (Callanan et al.\\ 1992; Pav\\-len\\-ko et al.\\ 1996; Belloni et al.\\ 1997; Steiman-Cameron \\& Scargle 1997; Iwasawa et al.\\ 1998) suggest that some periodic components are present in the signal. Thus we conclude with M.~J.\\ Rees (1998): {\\sl{``There is a real chance that someday observers will find evidence that an AGN is being modulated by an orbiting star, which would act as a test particle whose orbital precession would probe the metric in the domain where the distinctive features of the Kerr geometry should show up clearly.''}} \\def\\lb#1{{\\protect\\linebreak[#1]}}" + }, + "9802/astro-ph9802355_arXiv.txt": { + "abstract": "As a continuation of our study of the faint galaxy luminosity function in the Coma cluster of galaxies, we report here on the first spectroscopic observations of very faint galaxies (R $\\le$ 21.5) in the direction of the core of this cluster. Out of these 34 galaxies, only one may have a redshift consistent with Coma, all others are background objects. The predicted number of Coma galaxies is 6.7 $\\pm$ 6.0, according to Bernstein et al. (1995, B95). If we add the 17 galaxies observed by Secker (1997), we end up with 5 galaxies belonging to Coma, while the expected number is 16.0 $\\pm$ 11.0 according to B95. Notice that these two independent surveys lead to the same results. Although the observations and predicted values agree within the error, such results raise into question the validity of statistical subtraction of background objects commonly used to derive cluster luminosity functions (e.g. B95). More spectroscopic observations in this cluster and others are therefore urgently needed. As a by-product, we report on the discovery of a physical structure at a redshift z$\\simeq$0.5. ", + "introduction": "The Coma cluster of galaxies is one of the most intensively studied clusters in the sky. Its high richness and low redshift have led it to become the archetypal target for both observational and theoretical cluster studies. However, one aspect of Coma that has not yet been fully investigated is the distribution and dynamics of faint cluster galaxies. Indeed, we do not fully know what contribution galaxies fainter than the photographic plate limit (cf. the photometry by Godwin et al. 1983) make to the overall light (and mass) of Coma, and whether this contribution can significantly increase the known baryonic mass in the cluster. Futhermore, the exact shape of the Coma luminosity function (Bernstein et al. 1995, hereafter B95, Biviano et al. 1995, Lobo et al. 1997, Trentham 1997) remains in doubt, with suggestions varying from a single atypical Schechter function with a steep faint end slope to a Gaussian at bright magnitudes and a power-law at the faint end. This confusion comes from the reliance on a 2D statistical correction to subtract the background contribution. Most papers analyzing cluster luminosity functions, especially in distant clusters, use this method (see e.g. Driver et al. 1994, De Propris et al. 1995), but such a procedure could induce severe correlated errors and biases. Although the large error bars on these predictions do not totally prohibit a conclusive statement on potential errors in the 2D statistical background subtraction, it is crucial to check carefully this method. Coma is an excellent first candidate for this purpose. In the present paper, we assume H$_0$=100 km~s$^{-1}$Mpc$^{-1}$ and q$_0$=0. ", + "conclusions": "As presented above, 34 faint galaxy redshifts have been obtained in the direction of the Coma cluster core: 29 galaxies have magnitudes in the range 19$\\leq$R$\\leq$21.3, and 5 do not have measured magnitudes. Out of these 34 galaxies, only one may have a redshift consistent with the Coma cluster, all the others being background objects. If we add the S97 results, we obtain a more representative sample of 46 redshifts along the Coma line of sight, out of which only 5 are in the cluster. Notice that these two independent surveys lead to the same results. Although the number of galaxies found to be in Coma is consistent at the 1$\\sigma$ level with the expected number of objects in Coma predicted by B95 (16$\\pm$11), such a result obviously raises into question the validity of the statistical subtraction of background objects. Tidal disruption can be proposed as an explanation of the apparent deficit of faint galaxies in this region of the core of the Coma cluster. Moreover, the presence of a distant cluster on the same line of sight probably increases the field counts. However, more spectroscopic observations of very faint galaxies are obviously needed; we are pursuing more data in this Coma field, but prefer to release the present data to encourage other observers to also obtain faint redshifts in the direction of Coma or other clusters, with larger telescopes if possible." + }, + "9802/astro-ph9802163_arXiv.txt": { + "abstract": "I briefly review the current status of observations of AGN-powered UV/optical light, starlight, dust and outflow phenomena in high-redshift powerful radio galaxies. The existing data are consistent with the hypothesis that powerful radio galaxies undergo a major episode of star formation at high redshift ($z\\simgt 4$) during which they form most of their stars, and subsequently evolve `passively', with the UV continuum emission in the $z\\sim 1$ galaxies being dominated by AGN-related processes rather than starlight from the underlying, aging population. ", + "introduction": "The host galaxies of distant radio sources are of fundamental importance to studies of galaxy formation and evolution primarily because these objects are the best candidates for the progenitors of present-day massive galaxies and represent strongly biased peaks in the matter distribution. This hypothesis is supported at low redshift by the association of powerful radio sources with gE and cD galaxies (Matthews \\etal 1964), at intermediate and high redshifts by the tendency for these sources to reside in moderately rich cluster environments (\\eg, Hill \\& Lilly 1991, Dickinson 1997), and by a few direct kinematic measurements of the masses of high-redshift powerful radio galaxies (\\eg, 3C265: Dey \\& Spinrad 1996). In this contribution, I will briefly review the state of the observational data on high-redshift powerful radio galaxies (hereafter HzRGs) and their relevance to our understanding of the evolution of these systems. \\begin{figure}[t] \\plotfiddle{adey_fig1.ps}{2.8in}{-90}{47}{47}{-190}{250} \\caption{Spectropolarimetric observations of two radio galaxies: 3C368 at $z=1.132$ (left panel) and 3C441 at $z=0.707$ (right panel). Since 3C368 is unpolarized, with $P<3\\%$, only the Stokes parameters ($Q$ and $U$) and the unbiased percentage polarization ($P_{unb}$) are shown. In contrast, 3C441 is strongly polarized and shows a monotonically decreasing $P$ and wavelength-independent polarized flux. Note the CaII~K stellar absorption in 3C441, and the polarization of the MgII$\\lambda$2800 emission line.} \\label{3c441pol} \\end{figure} ", + "conclusions": "I have attempted to present an overview of our observations of $z>1$ radio galaxies and describe our present observational understanding of the evolution of the different components in these objects. The preliminary evidence suggests that the AGN contribution to the UV light is less important at $z\\sim 4$ than at $z\\sim 1$ and may reflect the spectral evolution of the stelar component which dominates the UV light at higher redshifts. Indeed, the present data, albeit sparse, qualitatively supports an evolutionary scenario in which powerful radio galaxies form the bulk of their stars before $z\\sim 3.5-4$, and then evolve relatively quiescently (\\ie, with little or no continuing star-formation) to $z\\simlt1$. This subject remains in its infancy and is still photon-starved. The spectroscopic and polarimetric observations that are necessary to elucidate the content and evolution of HzRGs require long exposures on the largest telescopes, and our understanding will therefore improve considerably during the next decade with the availability of sensitive instruments on the Keck, VLT and Gemini telescopes. \\begin{acknow} I thank my collaborators Wil van Breugel, Hy Spinrad, Bill Vacca, Daniel Stern, Ski Antonucci, Mark Dickinson, Andrea Cimatti, Andrew Bunker and Huub Rottgering for permitting me to present some of our results prior to publication. Much of the data presented here were obtained at the W.~M.~Keck Observatories, and I thank the Observatory staff for their expert assistance. I am very grateful to the KNAW and to NOAO for financial support that made it possible for me to attend the conference. In particular, I thank Huub R\\\"ottgering, Philip Best, Matt Lehnert and George Miley for providing me with the opportunity to visit Amsterdam, for a stimulating conference, and particularly for their extreme patience in waiting for this contribution. \\end{acknow}" + }, + "9802/astro-ph9802138_arXiv.txt": { + "abstract": "We investigate the effects of discontinuous mass loss in recurrent outburst events on the long--term evolution of cataclysmic variables (CVs). Similarly we consider the effects of frictional angular momentum loss (FAML), i.e.~interaction of the expanding nova envelope with the secondary. The Bondi--Hoyle accretion model is used to parameterize FAML in terms of the expansion velocity $v_{\\rm exp}$ of the nova envelope at the location of the secondary; we find that small $v_{\\rm exp}$ causes strong FAML. Numerical calculations of CV evolution over a wide range of parameters demonstrate the equivalence of a discontinuous sequence of nova cycles and the corresponding mean evolution (replacing envelope ejection by a continuous wind), even close to mass transfer instability. A formal stability analysis of discontinuous mass transfer confirms this, independent of details of the FAML model. FAML is a consequential angular momentum loss which amplifies the mass transfer rate driven by systemic angular momentum losses such as magnetic braking. We show that for a given $v_{\\rm exp}$ and white dwarf mass the amplification increases with secondary mass and is significant only close to the largest secondary mass consistent with mass transfer stability. The amplification factor is independent of the envelope mass ejected during the outburst, whereas the mass transfer amplitude induced by individual nova outbursts is proportional to it. In sequences calculated with nova model parameters taken from Prialnik \\& Kovetz \\cite{prialnik:kovetz} FAML amplification is negligible, but the outburst amplitude in systems below the period gap with a white dwarf mass $\\simeq 0.6 \\msun$ is larger than a factor of 10. The mass transfer rate in such systems is smaller than $10^{-11} \\msun$/yr for $\\simeq 0.5$~Myr ($\\simeq 10\\%$ of the nova cycle) after the outburst. This offers an explanation for intrinsically unusually faint CVs below the period gap. ", + "introduction": "Cataclysmic variables (CVs) are short--period binary systems in which a Roche--lobe filling low--mass main--sequence secondary transfers mass to a white dwarf (WD) primary. The transferred matter accretes onto the WD either through a disc or a stream and slowly builds up a hydrogen--rich surface layer on the WD. With continuing accretion the pressure at the bottom of this layer increases, and hydrogen burning eventually starts. The thermodynamic conditions at ignition determine how the burning proceeds (e.g.\\ Fujimoto 1982). If the degeneracy is very high, a thermonuclear runaway occurs, leading to a violent outburst terminated by the ejection of all or most of the accumulated envelope. Classical novae are thought to be objects undergoing such an outburst (cf.\\ Livio 1994 for a recent review). Ignition at moderate or weak degeneracy causes strong or weak H shell flashes, whereas stable or stationary hydrogen burning requires fairly high accretion rates $\\ga 10^{-7} \\msun/$yr which are not expected to occur in CVs. Mass transfer in CVs is driven by orbital angular momentum losses which generally shrink the binary and maintain the semi--detached state. The observed properties of short--period CVs below the CV period gap (orbital period $P \\la 2$~h) are consistent with gravitational wave radiation as the only driving mechanism. A much stronger angular momentum loss, usually assumed to be magnetic stellar wind braking, is needed for systems above the gap ($P \\ga 3$~h). The assumption that magnetic braking ceases to be effective once the secondary becomes fully convective in turn provides a natural explanation for the period gap as a period regime where the systems are detached and therefore unobservable (Spruit \\& Ritter 1983, Rappaport et al.\\ 1983). The resulting typical mass transfer rate $X$ in CVs is $X \\simeq 5 \\times 10^{-11} \\msun/$yr below the gap and $X \\simeq 10^{-9} - 10^{-8} \\msun/$yr above the gap (see e.g.\\ King 1988, Kolb 1996, for reviews). For negligible wind losses from the system the accretion rate is essentially the same as the transfer rate. With the above typical values the H ignition on the WD turns out to be degenerate enough to cause more or less violent outbursts (e.g.\\ Prialnik \\& Kovetz 1995). As the mass to be accumulated before ignition is very small ($\\mig \\simeq 10^{-6} - 10^{-3} \\msun$) the outbursts recur on a time $t_{\\rm rec} = \\mig/X$ much shorter than the mass transfer timescale which determines the long--term evolution. Studies of the secular evolution of CVs make use of this fact and replace a sequence of nova outbursts with given recurrence time $t_{\\rm rec}$ and ejected envelope mass $\\mej$ by a continuous isotropic wind loss from the WD at a constant rate $\\mej/t_{\\rm rec}$. Such a procedure obviously neglects any effect that nova outbursts may have on the long--term evolution. These are in particular \\\\ (i)~The evolution of the system is not continuous but characterized by sudden changes of the orbital parameters, causing the mass transfer rate to fluctuate around the continuous wind average value (see Sect.~2). It is not a priori clear if the continuous wind average properly describes the system's evolution close to mass transfer instability. \\\\ (ii)~At visual maximum (and the following decline) nova envelopes have pseudo--photospheric radii of typical giants, i.e.~much larger than the orbital separation. Therefore the secondary is engulfed in this envelope and possibly interacting with it. Drag forces on the secondary moving within the envelope can lead to frictional angular momentum loss (FAML) from the orbit, and accretion of envelope material onto the secondary could increase its photospheric metal abundances (pollution). \\\\ (iii)~The H burning hot WD is extremely luminous ($\\sim 10^{\\rm 4} L_{\\odot}$, compared with accretion luminosities $\\sim 1 L_{\\odot}$) for as long as a few years. This might drive additional mass loss from the secondary star. In this paper we will focus on the effects of mass loss discontinuities and FAML. We neglect irradiation as it does not last long enough to influence the long--term evolution. Pollution is expected to be important only for metal poor secondaries (Stehle 1993); we neglect it altogether. In Sect.~2 we formally derive the continuous wind average and consider the stability of mass transfer in the presence of nova discontinuities analytically, with FAML of arbitrary strength. We review previous studies on FAML and follow Livio et al.\\ \\cite{livio:etal} to derive a simple quantitative model for FAML in Sect.~3. Using this description we perform numerical calculations of the long--term evolution of CVs with various strengths of FAML, both for sequences of nova outbursts and the continuous wind average. Results of such computations where the FAML strength and the ejected mass per outburst have been varied systematically are shown in Sect.~4. Sequences with FAML parameters taken from the consistent set of nova models by Prialnik \\& Kovetz \\cite{prialnik:kovetz} are shown at the end of Sect.~4. Section 5 discusses our results. ", + "conclusions": "In this paper we considered the effects of nova outbursts on the secular evolution of CVs. As a result of these outbursts the secular mean mass transfer rate and the orbital period are not continuous functions of time but change essentially discontinuously with every nova outburst by an amount proportional to the ejected envelope mass. In addition, energy and angular momentum can be removed from the orbit due to dynamical friction of the secondary orbiting in the expanding nova envelope. The discontinuous evolution with a given strength of frictional angular momentum loss (FAML) is usually replaced by the corresponding continuous wind average evolution, where the mass and angular momentum loss associated with a nova outburst is assumed to be distributed over the inter--outburst time and to form an isotropic wind from the white dwarf. We showed analytically that the well--known mass transfer stability criterion for the latter case can also be derived from a proper analysis of the real, discontinuous process, for an arbitrary strength of FAML--amplification. We specified a quantitative model for FAML within the framework of Bondi--Hoyle accretion following Livio et al.\\ \\cite{livio:etal}. In this model the strength of FAML depends crucially on the expansion velocity $v_{\\rm exp}$ of the envelope at the location of the secondary, being the stronger the smaller $v_{\\rm exp}$ is. We expect that the resulting simple one--parameter description properly describes the order of magnitude of the FAML effect and, more importantly, the differential dependences on fundamental binary parameters. Hence although it is a useful way to study the potential influence of FAML systematically, it certainly cannot replace a detailed modelling of the frictional processes. Calculations of the long--term evolution of CVs verified the validity of the replacement of the discontinuous sequence of nova cycles with the continuous wind average, even for situations close to mass transfer instability, whatever the strength of FAML. The mass transfer rate in the continuous wind average evolution is FAML--amplified, i.e.\\ by the factor (\\ref{eq:mdotratio}) larger than the transfer rate driven by systemic angular momentum losses alone. For a given CV this factor is determined by $K_1$ alone, i.e.\\ independent of the ejection mass. We emphasize that FAML only {\\em amplifies} the transfer rate caused by systemic losses, it does not {\\em add} to them. Hence FAML is a particular example of consequential angular momentum losses (CAML) investigated in detail by King \\& Kolb \\cite{king:kolb}. In general, the FAML amplification factor turns out to be large only when the envelope expansion is very slow ($K_1 \\la 0.5$, i.e.\\ $v_{\\rm exp} \\la 200$~km/s) {\\em and} when the system is already close to thermal mass transfer instability (Fig.\\ref{fig:factor_fix}). This latter condition means that e.g.\\ for an evolution with strong FAML $K_1=0.1=$~const.\\ the averaged mass transfer rate is significantly affected only at long orbital period, $P \\ga 5$~h. The magnitude and direction of the outburst amplitudes of the mass transfer rate $X$ and the orbital period $P$ depend on both the ejection mass $\\mej$ and the FAML parameter $K_1$. For weak (or negligible) FAML the outbursts are towards lower mass transfer rate and longer orbital period, for strong FAML towards larger $X$ and shorter $P$. There is an intermediate regime where the outburst amplitudes essentially disappear. Theoretical models for nova outbursts generally find larger expansion velocities (e.g.\\ Prialnik 1986, Prialnik \\& Kovetz 1995), typically $K_1 \\ga 1$. This is certainly true for the terminal velocities, but in more recent models probably also for the crucial velocity at the secondary's location, i.e.\\ closer to the WD. Kato \\& Hachisu \\cite{kato:hachisu} argue that in the wind mass loss phase of a nova outburst the main acceleration (at about the sonic point) takes place at a temperature where the opacity has a maximum. Thus the introduction of the OPAL opacities had the effect of moving this point closer to the WD. A comparison of the radial velocity profiles found by Kato \\& Hachisu \\cite{kato:hachisu} with those obtained previously (e.g.~Prialnik 1986, Kato 1983) confirms this. As a result, in the newer models the velocities at the location of the secondary are already quite close to their terminal values. This explains why Kato \\& Hachisu \\cite{kato:hachisu} find only marginal effects from dynamical friction. This seems to suggest that the overall influence of FAML on the long--term evolution of CVs is small. However, in view of the considerable simplifications of our FAML description and the uncertainties of theoretical TNR models for nova outbursts, it is worthwhile to investigate sytematically mass transfer stability with FAML, and to consider the role of a feedback between outburst characteristics (hence FAML strength) and the mean mass transfer rate prior to the outburst. We will study this in a forthcoming paper (Kolb et al., in preparation). To illustrate further the effect FAML might have on the secular evolution of CVs we have calculated evolutionary sequences with envelope expansion velocities, ignition and ejection masses taken from the extended set of nova models by Prialnik \\& Kovetz \\cite{prialnik:kovetz}. As expected, the continuous wind average evolution hardly differs from the standard CV evolution without FAML. As a consequence of the large ejection mass (several $\\simeq 10^{-4} \\msun$/yr) the outburst amplitudes become very large (a factor $\\ga 10$) below the period gap for intermediate mass WDs ($M_1\\simeq 0.6\\msun$) --- even more so as the outburst--induced decrease of the mass transfer is largest if FAML vanishes. Such systems have a mass transfer rate less than $10^{-11} \\msun$/yr for $\\simeq 0.5$~Myr after the outburst, and could account for intrinsically faint CVs below the period gap. We speculated (Sect.~4.5) if TOADs (e.g.\\ Sproats et al.\\ 1996) could represent such systems. We finally note that the apparent scatter in observationally derived values for the mass transfer rate of CVs with comparable orbital period, well--known since the review of Patterson \\cite{patterson}, is unlikely to be due to outburst amplitudes. First, we expect from Fig.~\\ref{fig:dina_high} that these amplitudes are small or negligible above the period gap, and second the systems spend most of the time close to the continuous wind average mass transfer rate (cf. Fig.~\\ref{fig:problow_t}). A more promising explanation for this scatter assumes mass transfer cycles which could be irradiation--induced (e.g.\\ King et al.\\ 1995, 1996). \\subparagraph{Acknowledgements.} We thank D.~Prialnik and A.~Kovetz for giving access to nova model parameters prior to publication, and M.~Livio and L.~Yungelson for providing a copy of an unpublished manuscript on FAML. K.S.~would like to thank M.~Ruffert for many discussions on Bondi--Hoyle accretion. We thank A.~King for improving the language of the manuscript. K.S.~obtained partial financial support from the Swiss National Science Foundation. Theoretical astrophysics research at Leicester is supported by a PPARC rolling grant." + }, + "9802/astro-ph9802281_arXiv.txt": { + "abstract": "The evolution of QSO clustering is investigated with a new sample of 388 QSOs with $0.30.6$ as a lower limit.} \\label{fig:bigpicture} \\end{figure} ", + "conclusions": "We have examined the constraints in the $\\Omega_0$--$\\Omega_\\Lambda$ plane arising from a combination of SN-Ia and CMB data. Our most important result is that the two data sets provide approximately orthogonal constraints and thus nearly maximal complementarity. We have illustrated this by a likelihood analysis of the current data. Even though very different uncertainties affect the two data sets, the likelihood functions are compatible, with the allowed region shown in Fig.~\\ref{fig:both}. As more data is acquired, the combination will serve as an important cross check. Due to its current fashionability, and for contrast, we have also looked at constraints on universes with flat spatial hypersurfaces and low-$\\Omega_0$, with $1-\\Omega_0$ in a smooth component called $X$. As expected the constraints on this model are much weaker since -- by design -- the CMB peak location varies little with the model parameters. The allowed region is near $w\\approx -1$ and the CMB peak does not prefer any value of $\\Omega_0$ within this region. What about the future of this enterprise? Both the Supernova Cosmology Project and the High-Z team have $\\sim50$ supernovae still to be analyzed, which should shrink the contours in Fig.~\\ref{fig:sn1} considerably. Multicolour data will help control the uncertainty due to redenning and the allowed region should lie along a line of slope $\\simeq1$ (most SN-Ia will be at $z\\sim{1\\over2}$) in the $\\Omega_0$--$\\Omega_\\Lambda$ plane with width $\\pm0.1$. On the CMB front, data from currently operational experiments could determine the location of the peak in $\\ell$ to $\\Delta\\ell\\la 30$ within the next year (assuming the peak is near $\\ell\\simeq 250$). {}From Fig.~\\ref{fig:bigpicture} we see that such a measurement of $\\ell_{\\rm peak}$ would be comparable to, but orthogonal to, the supernova constraint. The first such experiment, the Mobile Anisotropy Telescope\\footnote{http://dept.physics.upenn.edu/$\\sim$www/astro-cosmo/ devlin/project.html} has already had a ``season'' in Chile. The Viper\\footnote{http://cmbr.phys.cmu.edu/vip.html} telescope is operating at the South Pole and results from several other experiments (see Bennett et al.~\\cite{PhysTod}; Table~1) are expected this year or next. And of course we anticipate that {\\sl MAP}, scheduled for launch in late 2000, will determine $\\ell_{\\rm peak}$ to a few per cent, the precise number depending on the angular scale of the peak. \\bigskip I would like to thank Joanne Cohn for a careful reading of the manuscript, Wayne Hu for useful conversations on ``generalized'' dark matter and scalar fields, Bob Kirshner and Saul Perlmutter for discussions on the supernova constraints, Douglas Scott for comments on the manuscript and Limin Wang for communicating his work on the cluster abundance ahead of publication." + }, + "9802/astro-ph9802126_arXiv.txt": { + "abstract": "We analyzed the ASCA X-ray data of 40 nearby clusters of galaxies, whose intracluster-medium temperature distributes in the range of 0.9--10 keV. We measured the Si and Fe abundances of the intracluster medium, spatially averaging over each cluster, but excluding the central $\\sim 0.15 h_{50}^{-1}$ Mpc region in order to avoid any possible abundance gradients and complex temperature structures. The Fe abundances of these clusters are 0.2--0.3 solar, with only weak dependence on the temperature of the intracluster medium, hence on the cluster richness. In contrast, the Si abundance is observed to increase from 0.3 to 0.6--0.7 solar from the poorer to richer clusters. These results suggest that the supernovae of both type-Ia and type-II significantly contribute to the metal enrichment of the intracluster medium, with the relative contribution of type-II supernovae increasing towards richer clusters. We suggest a possibility that a considerable fraction of type-II supernova products escaped from poorer systems. ", + "introduction": "The X-ray emitting hot intracluster medium (ICM) of clusters of galaxies is known to contain a large amount of heavy elements, presumably processed in the stellar interior and ejected into the intracluster space (e.g. Hatsukade 1989; Arnaud et al. 1992; Tsuru 1992). However, the mechanisms which supplied such a large amount of metals to intracluster space has not yet been understood well, due to insufficient knowledge of the abundance ratios and spatial metallicity distributions in the ICM. In particular, relative importance of type-Ia supernovae (SNe Ia; e.g. Mihara, Takahara 1994) and type-II supernovae (SNe II; e.g. Arnaud et al. 1992) has remained controversial. Since the SNe Ia products are iron-enriched, while the SNe II products are rich in $\\alpha$-elements, such as O, Ne, Mg, and Si, we vitally need to make separate measurements of the Fe abundance and those of $\\alpha$-elements. Using the X-ray data acquired with ASCA (Tanaka et al. 1994), Mushotzky et al. (1996) studied abundance ratios of four clusters with the ICM temperature of $kT$=3--4 keV, and found that $\\alpha$-elements in their ICM are roughly twice more abundant than Fe in solar abundance units. This was also confirmed by Tamura et al. (1996) and Xu et al. (1997). Mushotzky et al. (1996) hence suggest that SNe II are the dominant source of the metals in the ICM. However, a study of hotter (richer) and cooler (poorer) clusters remains to be performed. We present here a summary of ASCA measurements of the Fe and Si abundances in the ICM of 40 nearby clusters with various degrees of richness. The Hubble constant is expressed as $H_0$=50 $h_{50}$ km s$^{-1}$ Mpc$^{-1}$. The solar abundances refer to the solar photospheric values by Anders and Grevesse (1989), with (Fe/H)$_{\\odot}=4.68\\times10^{-5}$ and (Si/H)$_{\\odot}=3.55\\times10^{-5}$. ", + "conclusions": "We have measured spatially averaged Si and Fe abundances of the ICM of 40 nearby clusters, excluding the central regions. Our results concerning the Si abundance confirm the previous report by Mushotzky et al. (1996), and extend it to a much larger sample, in that the Si/Fe ratio is high at 1.5--2 (in solar unit) in clusters with considerable richness. This indicates that SNe II play a significant role in hotter clusters, because such abundance ratios differ from those found in the SNe Ia products (however see also Ishimaru, Arimoto 1997). We have also discovered that the Si/Fe abundance ratio decreases towards poorer clusters. If the metals are produced by the SNe II alone, this demands the chemical composition of the SNe II products to depend on the cluster richness, due, e.g., to the difference in the initial mass function of stars. This is, however, unlikely, since no difference has been seen between poorer and richer clusters in terms of the color-magnitude relation of member ellipticals (Visvanathan, Sandage 1977). Dust confinement of silicon in poor clusters is not likely either, because the dust evaporation time scale is at most $10^8$ yr in the ICM environment (Itoh 1989). Thus, the varying Si/Fe ratio may not be explained without invoking increasing contributions from SNe Ia in poorer clusters. This inference is reinforced by comparing clusters with supernova remnants (SNRs): the X-ray spectra of very poor ($kT\\sim1$ keV) clusters (e.g. Fukazawa et al. 1996) resemble very much those of type-Ia SNRs showing strong Fe-L lines, while distinct from those of type-II SNRs exhibiting prominent O-K, Ne-K, and Mg-K lines (Hayashi et al. 1994; Hughes et al. 1995). Furthermore, Ishimaru and Arimoto (1997) discuss that more than half the Fe in the ICM can be produced by SNe Ia, even in rich clusters. We therefore conclude that both types of supernovae contribute significantly to the metal enrichment of the ICM, with the relative contribution of SNe Ia increasing towards poorer clusters. This provides the first observational evidence that SNe Ia contribute significantly to the metal enrichment of the ICM. The richness-dependent relative contributions from the two types of SNe, in turn, may allow two alternative explanations. One is to invoke a richness dependence in their relative frequencies. This may be possible, because richer clusters clearly exhibit higher fractions of elliptical galaxies than do poorer clusters (Dressler 1980), and ellipticals are thought to have had higher frequencies of SN II per unit mass than spirals in the past (e.g. Arimoto, Yoshii 1987). However, the recent HST discovery of rapid morphological evolutions of member galaxies in rich clusters is against this possibility. (Dressler et al. 1994). Alternatively, the observed variation in the Si/Fe ratio may be explained if the SNe II products have higher specific energies than the SNe Ia products. Then, a larger fraction of SNe II products will escape from a shallower gravitational potential of poorer systems, while the SN Ia products may be confined even by the poorest systems, thus producing the observed trend in the Si/Fe ratio. This is quite plausible, because the SNe II products must have been supplied mostly in the form of galactic wind during early phase of cluster formation, whereas SNe Ia occurred on much prolonged time scales as isolated events. We expect this effect to have been even stronger during the galactic wind phase, when the gravitational potential may well have been shallower than today. Actually, there is good evidence that a large quantity of supernova products escaped from individual elliptical galaxies (Matsushita 1997), and to a less extent, from groups of galaxies (Fukazawa et al. 1996). Further investigation of the metal confinement efficiency will be presented in a forthcoming paper. \\vspace*{0.8cm} The authors are grateful to Prof. K. Yamashita and Prof. Y. Tanaka for valuable comments, and to the ASCA team for their help in the spacecraft operation and calibration. \\clearpage" + }, + "9802/astro-ph9802310_arXiv.txt": { + "abstract": "\\noindent We study the diffusive transport of neutrinos in a newly born neutron star to explore its sensitivity to dense matter properties. Energy and lepton number which are trapped during the catastrophic implosion diffuse out on the time scale of a few tens of seconds. Results for different dense matter models are presented. ", + "introduction": "\\vspace{-0.1in} The core of a massive star implodes when its mass exceeds the Chandrashekar mass. The hot and dense remnant formed subsequent to the implosion is a protoneutron star. Numerical simulations of the implosion (and the subsequent formation of a shock wave at core bounce) indicate that, due to the high densities and temperatures, most of the star's gravitational binding energy and lepton number released remains trapped in the star as neutrinos. The general features of the early evolution have been discussed in prior work \\cite{BL,KJ}. The object of this work is to elucidate the role played by the microphysical inputs (equation of state (EOS) and neutrino opacities) on the macrophysical evolution of the protoneutron star. \\vspace{-0.2in} ", + "conclusions": "" + }, + "9802/astro-ph9802195_arXiv.txt": { + "abstract": "It is shown that induced Raman scattering of electromagnetic waves in the strongly magnetized electron-positron plasma of pulsar magnetosphere may be important for wave propagation and as an effective saturation mechanism for electromagnetic instabilities. The frequencies, at which strong Raman scattering occurs in the outer parts of magnetosphere, fall into the observed radio band. The typical threshold intensities for the strong Raman scattering are of the order of the observed intensities, implying that pulsar magnetosphere may be optically thick to Raman scattering of electromagnetic waves. ", + "introduction": "We aim to estimate the effects of the the induced Raman scattering of the strong electromagnetic wave propagating in pulsar magnetosphere. Raman scattering may be considered as a parametric decay of the initial transverse electromagnetic wave in the another electromagnetic wave and plasma wave. The probability of this process is greatly enhanced if there many waves present in the final states. This is induced Raman scattering. Another type of induced scattering - induced Brullion scattering, i.e. the decay of the initial transverse electromagnetic wave in the another electromagnetic wave and ion sound is prohibited in electron-positron plasma, since pair plasma does not support low frequency, density perturbing waves (like ion sound wave in electron-ion plasma). Strong nonlinear coupling occurs when two waves beat together and the sum or difference frequency and wavelength match the frequency and wavelength of the third wave: \\begin{eqnarray} && \\omega_3= \\omega_1+ \\omega_2 \\mbox{} \\nonumber \\\\ \\mbox{} && k_{z,3}= k_{z,1}+ k_{z,2} \\label{n} \\end{eqnarray} In quantum language these relations may be interpreted as conservation of energy and momentum along the magnetic field, respectively. (the transverse components of the momentum are not conserved). The probability of the induced scattering depends on the number of waves in the final state, which, in turn, depends on the damping or escape rate of the daughter waves. The energy transfer between the modes will be efficient if the energy of the pump wave is strong enough to overcome the damping losses or escape of the generated waves. Thus, the induced Raman scattering is a threshold process: if the intensity of the pump wave exceeds the threshold value, the initial electromagnetic wave would start converting energy into the decay waves decreasing its amplitude exponentially. Induced Raman scattering may be important in the pulsar environment in two ways. First, it may provide en effective damping of the existing electromagnetic wave, that has been generated by some emission mechanism at the lower altitude in the pulsar magnetosphere, where the resonant conditions for the induced Raman scattering were not satisfied. This may result in a short time variability which is generally observed in pulsar radio emission. Secondly, it may provide an effective saturation mechanism for the growth of the electromagnetic wave provided that the conditions for the wave excitation by some mechanism are satisfied in the region where an effective Raman scattering takes place. The first possibility, i.e. the scattering of the existent wave, is simpler to consider. As a first approximation, we can treat the intensity of the pump wave as a constant. Then the nonlinear equations describing the wave coupling become linear in amplitudes of the decayed waves. The exponentially growing solutions will imply an effective energy transfer from the pump wave. As the intensities of the decay waves grow, this approximation breaks down in two cases: when the amplitudes of the decay waves become comparable to the pump wave or when the amplitude of the decay waves enter nonlinear stage and the waves start loosing energy due to some nonlinear process (like particle trapping and acceleration). The net effect of any energy loss by the decay wave is the depletion of the original pump wave. For the practical estimates, we can neglect the nonlinear stages of evolution (like cyclic energy transfer between the pump and the daughter waves) and assume that if the intensity of the pump wave begins to decrease exponentially in the linear stage, then the wave decays completely. The second possibility, e.g. when the induced Raman scattering provides a nonlinear saturation mechanism, is more complicated, since the intensities of all waves may be of the same order. A considerable simplification in this case may be obtained if the damping of one of the decay modes is very strong or if it leaves the region of the resonant interaction fast enough. Then, after a short period of time (when the intensities of the other weakly damped waves grow considerably) the intensity of this damped mode is much smaller than the intensities of the two other modes and can be neglected. In what follows we neglect the possible nonlinear stages of the development of Langmuir turbulence (particle trapping or quasilinear diffusion). We also assume that the pump wave is broadband. This is quite different from the conventional laboratory case of a monochromatic laser-plasma interaction. The condition of a broadband pump wave implies that its band width $\\Delta \\omega$ is much larger than the typical growth rate of the decay instability $\\Gamma$: \\begin{equation} \\Delta \\omega \\gg \\Gamma \\label{kd941} \\end{equation} If this condition is satisfied, then we can use a random phase approximation for the statistical description of the interacting waves. On the other hand, the condition of weak turbulence allows one to calculate the matrix elements of the interaction in the approximation of stationary phases. The general expression for the third order nonlinear current in plasma in a magnetic field has been written down by \\cite{TsytovichShvartsburg}. An extremely complicated form of the corresponding expressions makes the general case of Raman scattering very difficult to consider. Several important simplifications can be done when considering induced Raman scattering in the pulsar magnetosphere. First, the superstrong magnetic field allows an expansion of the currents in $1/\\omega_B$. Second, in the pair plasma with the same distributions of electrons and positrons some of the nonlinear currents cancel out since they are proportional to the third power of the electric charge (this cancellation is exact in the unmagnetized electron-positron plasma). The third, less justified approximation, is that we will make is that all the three interacting waves propagate along magnetic field. This is an important assumption. It allows us to simplify the consideration considerably and to obtain some analytical estimates of the effects of induced Raman scattering. A short overview of the work on the nonlinear process in the pulsar magnetosphere will be appropriate here. The possible decay processes for the transverse waves and the corresponding references are ($t$ denotes transverse wave, $l$ denotes longitudinal wave and $e$ denotes a charged particle): (i) a decay of a transverse wave into two transverse wave $ t \\rightarrow t ^{\\prime} + t ^{\\prime \\prime}$ (\\cite{GedalinMachabeliPismaAj}), (ii) a decay of a transverse wave into two Lagmiur waves $ t \\rightarrow l+l^{\\prime}$ (\\cite{GedalinMachabeli1983}), (iii) a decay of a transverse wave into another transverse and Lagmuir wave $ t \\rightarrow t ^{\\prime} + l$ (\\cite{GedalinMachabeli1983}) \\footnote{ The final answer for the matrix element contains probably an insignificant typographical error} (iv) induced scattering of transverse waves $ t + e \\rightarrow t ^{\\prime} + e $ (\\cite{BlandfordScharleman}, \\cite{SincellKrolik1992}, \\cite{OchelkovUsov}, \\cite{Wilson82}). The possible processes for Lagmuir waves are (i) a decay of a Lagmuir wave into another Lagmuir and transverse wave $ l \\rightarrow t + l ^{\\prime}$ (\\cite{GedalinMachabeli1983}), (ii) a decay of a Lagmuir wave into two transverse waves $ l \\rightarrow t + t ^{\\prime}$ (\\cite{Mikhailovski1980}), (iii) Lagmuir wave merger $ l + l ^{\\prime} \\rightarrow t$ (\\cite{MachabeliMamradzeMelikidze1982}), (iii) induced scattering of Lagmuir waves (\\cite{Lyubarskii93}). When treating the nonlinear process involving transverse waves, the matrix elements in above works have been calculated in the drift approximation - the expansion in parameter $1/\\omega_B$ was done in the very beginning ($\\omega_B= |q| B/m c$ is the positive nonrelativistic cyclotron frequency, $B$ is magnetic field, $m$ is mass of an electron and $c$ is the speed of light). ", + "conclusions": "\\label{Conclus} In this work we considered induced Raman scattering of a transverse electromagnetic wave propagating along magnetic field in a pair plasma. We found that induced Raman scattering can be an important factor for the wave propagation in the pulsar magnetosphere and as a nonlinear saturation mechanism for the electromagnetic instabilities ( cyclotron-Cherenkov instability, for example). The threshold intensities for strong Raman scattering, which are determined by the damping rate of the Langmuir waves in the relativistic plasma, are comparable to the observed fluxes from radio pulsars. One of the major limitations of this work is the assumption of the propagation along the magnetic field. The natural extension of this work to oblique propagation would involve several complications: (i) there will be a distinction in the transition probabilities between extraordinary mode (with the electric field of the waves perpendicular to the ${\\bf k} - {\\bf B}$ plane) and ordinary mode (with the electric field in the ${\\bf k} - {\\bf B}$ plane), (ii) ordinary mode has an sensitive dependence on the angle of propagation, so that for the ordinary mode the Raman scattering will occur only in a very limited angle around magnetic field; the scattered waves will be a backward propagating, (iii) for the extraordinary mode the Raman scattering is possible for large angles of propagation; in this case both electromagnetic waves can be forward propagating." + }, + "9802/nucl-th9802004_arXiv.txt": { + "abstract": "The $^3{\\rm He}({^3{\\rm He}},2p){^4{\\rm He}}$ and $^3{\\rm H}({^3{\\rm H}},2n){^4{\\rm He}}$ reactions are studied in a microscopic cluster model. We search for resonances in the $^3$He+$^3$He and $^4{\\rm He}+p+p$ channels using methods that treat the two- and three-body resonance asymptotics correctly. Our results show that the existence of a low-energy resonance or virtual state, which could influence the $^7$Be and $^8$B solar neutrino fluxes, is rather unlikely. Our calculated $^3{\\rm He}({^3{\\rm He}},2p){^4{\\rm He}}$ and $^3{\\rm H}({^3{\\rm H}},2n){^4{\\rm He}}$ cross sections are in a good general agreement with the experimental data. ", + "introduction": "One way to test models of the solar interior is by observing the neutrinos generated by the nuclear reaction network which is the solar energy source \\cite{SNexp}. As a striking and exciting fact, all terrestrial solar neutrino experiments observe fewer neutrinos than predicted by standard solar models \\cite{SNtheo}. The picture which emerges from the various experiments is \\cite{Hata} that the $^8$B flux is about half of its predicted value ($\\phi_8=0.4\\phi_8^{SSM}$), while the $^7$Be neutrinos appear to be completely missing ($\\phi_7=0$). Here SSM refers to the standard solar model of Bahcall \\cite{SNtheo}. The $^7$Be nuclei, which are the seeds of both the $^7$Be and $^8$B neutrinos, are produced in the $^4{\\rm He}({^3{\\rm He}},\\gamma){^7{\\rm Be}}$ reaction. This reaction competes with the $^3{\\rm He}({^3{\\rm He}},2p){^4{ \\rm He}}$ process, which is the final step of the first branch of solar hydrogen burning (ppI chain) \\cite{Bahcall}. If the cross section of the $^3{\\rm He}({^3{\\rm He}},2p){^4{ \\rm He}}$ reaction were larger than believed then the probability of the $^4{\\rm He}({^3{\\rm He}},\\gamma){^7{\\rm Be}}$ branch would be smaller, and hence the $\\phi_8$ and $\\phi_7$ fluxes would be suppressed (without, however, significantly affecting the unexpected $\\phi_7 / \\phi_8$ ratio deduced from the neutrino experiments \\cite{Raghavan}). To increase the $^3{\\rm He}({^3{\\rm He}},2p){^4{\\rm He}}$ reaction rate, a hypothetical resonance at low energies had been suggested \\cite{Fowler} and looked for in various experiments (see \\cite{LUNA} for the most recent work). The importance of the $^3{\\rm He}({^3{\\rm He}},2p){^4{\\rm He}}$ reaction has led to continued experimental efforts to measure the cross section down to solar energies. In the latest experiment \\cite{LUNA} the LUNA collaboration measured the cross section down to $E_{cm}=20.76$ keV, which is well within the region of the most effective solar energies (Gamow window). Although, they do not see any evidence of a possible resonance, the existence of such a state at still lower energies cannot be {\\it a priori} ruled out yet. In the present work we study the $^3{\\rm He}({^3{ \\rm He}},2p){^4{\\rm He}}$ and the mirror $^3{\\rm H}({^3{\\rm H}},2n){^4{\\rm He}}$ reactions in a microscopic cluster model. We search for signs of possible resonances and study the energy dependence of the reaction cross sections. ", + "conclusions": "We have studied the $^3{\\rm He}({^3{\\rm He}},2p){^4{\\rm He}}$ and $^3{\\rm H}({^3{\\rm H}},2n){^4{\\rm He}}$ reactions within the microscopic cluster model using significantly larger model spaces than previously employed. Our motivation and results have been twofold: We searched for signs of possible low-energy resonances in $^3{\\rm He}({^3{\\rm He}},2p){^4{\\rm He}}$, which could have important effects on the $^7$Be and $^8$B solar neutrino fluxes. The $^3{\\rm He}+{^3{\\rm He}}$ and $^4{\\rm He}+p+p$ channels were studied separately, which allowed us to use methods that can treat the two- and three-body asymptotics in a rigorous way. We extended the two- and three-body scattering matrices to complex energies and searched for their poles. We have not found any indication for the existence of a low-energy resonance or virtual state that could cause observable effects in the cross section. Thus, it is unlikely that a yet unobserved resonance around the threshold energy in the $^3$He($^3$He,2p)$^4$He reaction might affect this important solar reaction cross section. We calculated the cross sections of the $^3{\\rm He}({^3{\\rm He}},2p){^4{\\rm He}}$ and $^3{\\rm H}({^3{\\rm H}},2n){^4{\\rm He}}$ reactions in the continuum-discretized coupled channel approximation. Our results are in a good general agreement with available data, except for the very low-energy $^3{\\rm H}({^3{\\rm H}},2n){^4{\\rm He}}$ cross section, where we observe a systematic deviation from the most precise measurement. Our test calculations show that the energy dependence of the cross sections is hardly influenced by the details of the continuum discretization, but might be caused by the approximate treatment of the 3-body continuum. Here improvements are certainly warranted." + }, + "9802/astro-ph9802148_arXiv.txt": { + "abstract": "Although radio-quiet quasars (RQQs) constitute $\\stackrel{>}{\\sim}90\\%$ of optically-identified quasar samples their radio properties are only poorly understood. In this paper we present the results of a multi-frequency VLA study of 27 low-redshift RQQs. We detect radio emission from 20 objects, half of which are unresolved ($\\leq 0.24''$). In cases where significant structure can be resolved, double, triple and linear radio sources on scales of a few kpc are found. The radio emission (typically) has a steep spectrum ($\\alpha \\sim 0.7$, where $S \\propto \\nu^{-\\alpha}$), and high brightness temperatures ($T_{B} \\geq 10^{5}$~K) are measured in some of the radio components. The RQQs form a natural extension to the radio luminosity - absolute magnitude distribution of nearby Seyfert 1s. We conclude that a significant fraction of the radio emission in RQQs originates in a compact nuclear source directly associated with the quasar. There are no significant differences between the radio properties of RQQs with elliptical hosts and those in disc galaxies within the current sample. ", + "introduction": "Shortly after their initial identification in 1963 by Schmidt it became clear that less than 10\\% of quasars are strong sources of radio emission. The gap in the quasar radio luminosity function is extremely pronounced, with only a handful of objects occupying the region between quasars which are radio-loud and those which are radio-quiet (Kellermann et al. 1989; Miller, Peacock \\& Mead 1990; Miller, Rawlings \\& Saunders 1993, henceforth MRS93), and it is now widely accepted that there are two distinct populations of quasar. However, from $\\sim100~\\mu$m through to soft X-ray wavelengths the properties of radio-loud and radio-quiet quasars (RLQs and RQQs respectively) are generally very similar, to the extent that to first order an RLQ spectrum can be considered to be the spectrum of an RQQ with the addition of a strong power-law component in the radio. The presence or absence of this radio component must be a fundamental indicator of the processes occurring in the quasar, but there is still no consensus as to exactly what it signifies. Studies of less luminous, nearby active galaxies have shown that the radio-loud objects ({\\it ie} Radio Galaxies) are invariably elliptical systems whereas the radio-quiets (Seyferts) tend to be spirals, suggesting that some property of gas-rich disc galaxies inhibits the formation of large, powerful radio sources (see Osterbrock 1991). Two factors have encouraged the extension of this result to the higher redshifts and larger nuclear luminosities typical of quasars: the success of Unified Schemes which link Radio Galaxies and RLQs via beaming effects and viewing angle (see Urry \\& Padovani 1995) and the established fact that Seyfert 1 nuclei and RQQs form a continuous sequence in terms of their optical luminosities and have identical emission line characteristics (see Osterbrock 1991, Antonucci 1993). However, recent advances in ground- and space-based observing have allowed the host galaxies of nearby quasars to be imaged and {\\it reliably} classified for the first time. The results have shown that {\\it not all} RQQs lie in disc systems but that as many as 50\\% might be found in elliptical hosts ({\\it eg} Taylor et al. 1996, Disney et al. 1995, Lacy, Rawlings \\& Hill 1992, V\\'{e}ron-Cetty \\& Woltjer 1990). Indeed, there is some evidence that elliptical galaxies might account for all of the most optically luminous RQQs (Taylor et al. 1996). Clearly the simple `radio-loud $\\equiv$ elliptical, radio-quiet $\\equiv$ disc' picture can no longer be supported, and it has become more important than ever to determine in what respects radio-quiet quasars are different from their radio-loud counterparts. Unfortunately, the most obvious wavelength regime in which they differ - the radio - is also the regime in which least is known about the properties of RQQs. Most radio surveys of optically-selected quasar samples have lacked the sensitivity to detect the radio-quiet objects, and the few high-sensitivity surveys have tended to give only fluxes at a single frequency, with only limited information on radio structure and spectral index. Like Seyferts, RQQs are not radio silent and with sensitivities at the milliJansky level modern multi-element synthesis telescopes can easily detect the radio emission from many of the nearer ($z \\leq 0.3$) objects. In this paper we present the results of a multi-frequency, high-resolution radio survey of low-redshift ($0.03 \\leq z \\leq 0.3$), low-luminosity (M$_{V}>-26$) RQQs using the Very Large Array (VLA) in its high-resolution A-configuration. The aim of this work was to search for compact, non-thermal radio sources associated with the quasars, to determine their properties and to compare them with those of radio-loud objects. As well as our own observations, we also make use of data from the literature. Our RQQ sample comprises 27 objects, 19 of which have already been observed with A-array at 4.8~GHz by Kellermann et al. (1989) as part of their radio study of the Bright Quasar Survey (BQS; Schmidt \\& Green 1983, Green, Schmidt \\& Liebert 1986). We have observed all 27 RQQs at 1.4~GHz and completed the coverage of the sample at 4.8~GHz with observations of the 8 non-BQS quasars. In addition we obtained 8.4-GHz maps for 21 objects. Seventeen of the RQQs in our sample were also included in a major investigation into quasar host galaxies (Dunlop et al. 1993, Taylor et al. 1996). We use these objects to investigate the relationship between the host and the `radio loudness' of the quasar and refer to them as our `host galaxy subsample'. Two objects (PG~0007$+$106 \\& 1635$+$119) have unusually large radio luminosities and arguably are not true {\\it radio-quiet} quasars but belong to a separate class of their own. We use the term `radio-intermediate quasar' (RIQ) to describe these. The paper is structured as follows. Data acquisition and reduction are discussed in Section 2, followed by a description of the resulting radio images in Section 3. Section 4 of the paper summarises the debate over the origin of non-thermal radio emission in RQQs and Section 5 explores the implications of the current observations, and compares the RQQ emission with the radio properties of other types of active galactic nuclei (AGN). In Section 6 we examine the relationship between the radio properties of the 17 RQQs in our host galaxy subsample and the properties of the hosts themselves, and our conclusions are summarised in Section 7. We assume $H_{0} =50$~kms$^{-1}$Mpc$^{-1}$ and $\\Omega_{0} = 1$ throughout. ", + "conclusions": "We have measured the radio fluxes and spectral indices of a sample of 27 RQQs using the VLA in A-configuration. The angular resolution of the VLA is not sufficient to unambiguously determine the origin of the radio emission in these objects, but from the data we are able to conclude the following: {\\bf (1)} In 74\\% of the sample we detect a radio source which is coincident with the optical quasar to within the positional uncertainty of the optical measurements. {\\bf (2)} The spectral indices of the RQQs are generally steep ($\\alpha \\sim 0.7$, where $S \\propto \\nu^{-\\alpha}$), although two objects, which also exhibit variability and unusually large radio luminosities, have flat radio spectra. {\\bf (3)} Lower limits on the brightness temperatures of many of the radio sources place them at the upper end of the range expected for emission related to stellar processes ($T_{B} \\sim 10^{5}$~K) ({\\it ie} in the regime in which we would expect to see some spectral flattening if the emission was indeed stellar in origin). In some objects the brightness temperature is almost certainly many orders of magnitude greater than this, in which case the radio emission {\\it cannot} be produced by a starburst. Large numbers of supernovae occurring within a very small volume would be required in order to reproduce the observed radio luminosities. Future observations with greater angular resolution will place more rigorous constraints on all these parameters. {\\bf (4)} In nine (possibly ten) objects we are able to resolve radio structure, which takes the form of double, triple and linear sources on scales of a few kiloparsecs. These structures are consistent with the 100-pc-scale radio jets observed in low-redshift Seyfert nuclei, and constitute strong evidence that collimated ejection of radio plasma from the central engine is occurring in RQQs. {\\bf (5)} The distribution of radio luminosities in RQQs form a natural extension to those of Seyfert 1 nuclei. There appears to be a correlation between radio luminosity and the optical absolute magnitude of the quasar, suggesting a close relationship between the central engine and the mechanism responsible for the bulk of the radio emission. {\\bf (6)} There are no statistically significant differences in the present data between the radio properties of the RQQs in disc-dominated galaxies and those in elliptical hosts. Thus, our results are consistent with the conclusions of other recent radio studies of RQQs: that in at least some RQQs a significant contribution to the overall radio emission comes from a compact nuclear source which is directly associated with the central engine of the quasar, and which is qualitatively similar to the more powerful radio sources observed in RLQs." + }, + "9802/astro-ph9802181_arXiv.txt": { + "abstract": "We present a study of the properties of {\\em dusty} warm absorbers and point to some important consequences for the resulting X-ray absorption spectra of AGN. Pronounced effects of the presence of {\\em dust} in {\\em warm} material are (i) an apparent `flattening' of the observed X-ray spectrum in the \\ros band due to a sequence of absorption edges and a shift to lower gas ionization, (ii) the presence of a strong carbon edge at 0.28 keV, and (iii) the expectation of increased time variability of the warm absorber parameters. The first two effects can be drastic and may completely hamper an X-ray spectral fit with a dusty warm absorber even if a dust-{\\em free} one was successfully applied to the data. In order to demonstrate facets of the dusty warm absorber model and test the recently reported important, albeit indirect, evidence for dusty warm material in the Narrow-Line Seyfert 1 galaxy IRAS 17020+4544 we have analyzed \\ros PSPC and HRI observations of this galaxy. The X-ray spectrum can be successfully described by a single powerlaw with index \\G=--2.4 plus small excess cold absorption, or alternatively by a steeper intrinsic powerlaw (\\G $\\simeq$ --2.8) absorbed by a {\\em dusty} warm absorber. The findings are discussed in light of the NLSy1 character of IRAS 17020 and consequences for NLSy1s in general are pointed out. In particular, the presence of dusty warm gas results in a steeper intrinsic powerlaw than observed, thus exaggerating the `FeII problem'. It also implies weaker potential warm-absorber contribution to high-ionization Fe coronal lines. ", + "introduction": "Warm absorbers reveal their presence by imprinting absorption edges on the soft X-ray spectra of active galaxies (AGN). Many were found and studied on the basis of observations by \\ros (e.g., Nandra \\& Pounds 1992, Komossa 1997) and \\asca (e.g., Mihara et al. 1994, Reynolds 1997). Within the last year evidence has accumulated that some warm absorbers contain significant amounts of dust. This possibility was first suggested by Brandt et al. (1996) to explain the lack of excess X-ray {\\em cold} absorption despite strong optical reddening of the quasar IRAS 13349+2438. Models that explicitly include the dust-gas and dust-radiation interaction using the photoionization code {\\em Cloudy} (Ferland 1993) were calculated by Komossa \\& Fink (KoFi hereafter; 1996, 1997a-d). Depending on the warm absorber parameters, the presence of dust turned out to have a strong influence on the X-ray absorption structure. The models were applied to several Seyfert galaxies. NGC 3227 (KoFi 1996, 1997b) and NGC 3786 (KoFi 1997c) were shown to be very good candidates for dusty warm absorbers as judged from optical-UV reddening properties as well as {\\em successful X-ray spectral fits}. On the other hand, the bulk of the warm material in NGC 4051 was found to be dust-free (KoFi 1997a). Recently, Reynolds et al. (1997) also presented a {\\em Cloudy}-based dusty absorber model which they successfully applied to the outer warm absorber in MCG 6-30-15. Indirect evidence for the association of some warm absorbers with dust was discussed in Reynolds (1997). IRAS 17020, which is part of the present study, is a Narrow Line Seyfert\\,1 galaxy (NLSy1; Moran et al. 1996, Wisotzki \\& Bade 1997) at redshift $z$=0.06 (de Grijp et al. 1992). Wisotzki \\& Bade stressed the heavy reddening of the optical spectrum. They also presented a powerlaw fit to the \\ros PSPC spectrum and pointed to the discrepancy between the cold column density derived from optical reddening and the one from the X-ray fit. Leighly et al. (1997; L97) detected an oxygen OVII edge in the \\asca spectrum of IRAS 17020 as well as high optical polarization. They also confirmed the strong optical reddening and derived a corresponding gaseous column density of $N_{\\rm opt}$ = 4 $\\times~10^{21}$ cm$^{-2}$. On this basis they suggested the presence of a warm absorber with internal dust in IRAS 17020. To test whether the model of a warm absorber that {\\em includes} the presence of dust consistently fits the X-ray spectrum of IRAS 17020, we apply such a model to the (archival) \\ros PSPC spectrum of this source (Sect. 2). Motivated by increasing indirect evidence for the presence of dusty warm absorbers, in the main part (Sect. 3) we point out some important general properties of the dusty material that can result in strong modifications of the X-ray absorption structure as compared to the dust-{\\em free} case, and study consequences. Since the presence of dust influences the X-ray spectrum most strongly at soft energies, one signature being a pronounced carbon edge at 0.28 keV outside the \\asca sensitivity range, \\ros (Tr\\\"umper 1983) is particularly well suited for such a study. To search for time variability of IRAS 17020, and check for the potential contribution to the X-ray spectral complexity of this source from the closeby optically bright star SAO 46462, we also present (PI) HRI observations. ", + "conclusions": "\\subsection{Properties of dusty warm absorbers} Some general properties of dusty warm absorbers are visualized in Fig. \\ref{dust_seq}. The presence of dust modifies the X-ray absorption structure in two ways: \\\\ \\begin{figure*}[t] \\vbox{\\psfig{figure=L3540_4.ps,width=8.5cm,% bbllx=2.9cm,bblly=1.1cm,bburx=18.3cm,bbury=12.2cm,clip=}}\\par \\vspace*{-6.1cm}\\hspace*{8.7cm} \\vbox{\\psfig{figure=L3540_5.ps,width=8.5cm,% bbllx=2.9cm,bblly=1.1cm,bburx=18.3cm,bbury=12.2cm,clip=}}\\par \\vspace*{-0.3cm} \\caption[dust_seq]{Influence of dust mixed with the warm absorber on the X-ray absorption structure for two values of the warm column density ($\\log N_{\\rm w}$ = 21.6 and 22.4). In {\\bf (a)} the straight solid line corresponds to the unabsorbed intrinsic spectrum, the heavy solid line to a dust-free warm absorber (WA) with solar abundances, the dotted line to a dust-free WA with the same depleted abundances as in the dusty model, and the long-dashed line to the change in absorption structure after adding dust to the WA. In {\\bf (b)} (in addition to the larger $N_{\\rm w}$ a higher $U$ was chosen for presentation purposes) the heavy solid line again corresponds to a dust-free WA with solar abundances, the long-dashed line to a dusty WA (note that, for better presentation, the dust was even depleted by a factor $f_{\\rm d}$ = 0.5), and the short-dashed line to the same dusty WA but without the graphite species of dust. Clear changes in the absorption structure are revealed for the models that include dust. The abscissa brackets the \\ros energy range (0.1-2.4 keV); some prominent edges are marked. } \\label{dust_seq} \\end{figure*} (i) Gas-dust interactions influence the thermal conditions in the gas and change its ionization state (e.g. Draine \\& Salpeter 1979). In particular, for high ionization parameters dust effectively competes with the gas in absorbing photons (Laor \\& Draine 1993). Adding dust to large column density warm material (cf. Fig. \\ref{dust_seq}b) results in strongly increased soft X-ray absorption, partially due to the stronger temperature gradient across the absorber, with more gas in a cold state. \\\\ (ii) Dust scatters and absorbs the incident radiation (e.g. Martin 1970) and X-ray absorption edges are created by inner-shell photoionization of metals bound in the dust (cf. Figs 4,5 of Martin \\& Rouleau 1991). Strong edges are those of neutral carbon and oxygen (Fig. \\ref{dust_seq}a,b). The shift in edge energies due to chemical and solid state effects is only of the order of a few eV (Greaves et al. 1984). The modifications of the observed X-ray spectrum as compared to the dust-free case can be drastic and are important to take into account when interpreting the observed X-ray properties of AGN. Consequences of the presence of dust are (A) `Flattening effect': The sequence of individual edges (dust-ones adding to the gas-ones) and the shift to lower ionization lead to a `smoothing' and apparent `flattening' of the observed X-ray spectrum (Fig. \\ref{dust_seq}a). Individual edges are less pronounced as in the dust-free case (where OVII and/or OVIII often dominate). This `flattening effect' of dust (in fact opposite to that of a dust-free warm absorber, which effectively steepens the X-ray spectrum in the \\ros band) has consequences for NLSy1s in general (Sect. 3.3). Another consequence is the possibility to `hide' dusty warm absorbers in the \\ros spectra (as shown by the successful fit to IRAS 17020) if the column densities are not too large{\\footnote {\\small The possibility of the presence of dusty warm material in a sample of NLSy1s to explain their polarization properties is suggested in Grupe et al. (1998).}}, although effect (B) has to be taken care of (or else non-standard dust properties allowed; dust in other galaxies may be different from Galactic one, but standard dust properties are also usually assumed to estimate optical reddening). (B) Carbon edge: Often, the strongest dust-created edge is that of carbon at 0.28 keV, stemming from the graphite species of dust. Its presence may prevent a successful spectral fit and it is important to actually apply the model of a dusty warm absorber to the data. Whereas dust including the graphite species is well consistent with the X-ray spectra of NGC 3227 (KoFi 1997b) and NGC 3786 (KoFi 1997c) and pure graphite was favoured by Reynolds et al. (1997) for MCG 6-30-15, it hampered a successful fit in other cases. (C) Increased sensitivity to radiation pressure: Material of high ionization parameter and particularly dust particles are strongly subject to radiation pressure (e.g., Laor \\& Draine 1993, Chang et al. 1987, Binette et al. 1997), causing the gas to outflow (if not otherwise confined) and leading to temporal changes of the absorber properties. Indeed, strong variability in X-ray count rate (of factors 10-15) has been recently detected for the dusty-warm-absorber candidates NGC 3227 and NGC 3786 which could be explained by variability in column density or ionization state of the warm absorber (KoFi 1997b,c). \\subsection{The warm absorber in IRAS 17020} Given the potentially strong modifications of the X-ray spectrum in the presence of {\\em dusty} warm material, the successful X-ray spectral fit lends further support to the suggestion (L97) of the presence of a dusty warm absorber in IRAS 17020. Although there is evidence for small excess cold absorption (e.g. Fig. 1) no description of the soft X-ray spectrum that involves the {\\em large cold} column inferred from optical observations{\\footnote{\\small Just to recall the assumptions that go into the approach of relating optical reddening by dust with X-ray absorption by cold gas: (i) optical and X-ray continuum reach the observer along the same path, (ii) the extinction is not strongly variable with time, and (iii) the dust/gas ratio is approximately the Galactic value; in modelling the warm absorber we assume (iv) its ionization state to be dominated by photoionization and {\\em Cloudy} to be applicable, and again (v) dust properties similar to Galactic dust.}} could be found. The successful fit of a dusty warm absorber requires a steep underlying powerlaw of index \\G $\\approx -2.8$ (the data are consistent with a flatter slope, if the graphite component of dust is excluded). \\subsection{NLSy1 nature of IRAS 17020} NLSy1s like IRAS 17020 are generally characterized by narrow broad lines, strong FeII emission and steep X-ray spectra (see Brandt et al. 1997 for a recent discussion). Several models were suggested to explain their apparently steep soft X-ray spectra, like a strong soft excess on top of a flat powerlaw (e.g. Puchnarewicz et al. 1992) or a (dust-free) warm-absorbed flat powerlaw (e.g. KoFi 1997d). Recent observations reveal increasing X-ray spectral complexity, with often both, a warm absorber and a soft excess present. The presence of {\\em dusty} warm material adds further to this complexity. In particular, the effective flattening of the observed spectrum implies a steeper intrinsic one, thus exaggerating the problem of producing the strong FeII emission (that in standard models requires hard X-ray photons) in NLSy1s. The possibility of a warm-absorber contribution to the optical high-ionization lines in Seyferts and NLSy1s was studied by KoFi (e.g. 1997d, Fig. 6). Here, we note that in case the absorber is {\\em dusty}, weaker contribution to line emission is found. In particular, due to the binding of Fe in dust, negligible contribution to high-ionization Fe coronal lines is expected. Although we find hints for X-ray variability of IRAS 17020 of the order of 25\\% on timescales from several hundred seconds to years, this amplitude is remarkably small as compared to some other reported cases of Seyfert and NLSy1 variability. \\subsection{Concluding remarks} We have shown how the presence of dust in warm absorbers can lead to strong modifications of the observed X-ray spectrum. The model of a dusty warm absorber was successfully fit to the X-ray spectrum of the NLSy1 galaxy IRAS 17020. This corroborates the suggestion of the presence of dusty warm material in this galaxy made on the basis of optical properties (Wisotzki \\& Bade 1997, L97) and the detection of an oxygen edge in the \\asca spectrum (L97). This first good case for a {\\em dusty} warm absorber in a NLSy1-type galaxy, together with the other good candidates in the Sy\\,1.8 galaxy NGC 3786 (KoFi 1997c), the Sy\\,1.5 NGC 3227 (KoFi 1996,1997b), the Sy\\,1 MCG 6-30-15 (Reynolds et al. 1997), and the quasar IRAS 13349 (Brandt et al. 1996) suggests this component to be common in all types of AGN. Given the strong X-ray absorption edges of neutral dust-bound metals, dusty warm absorbers will play an important role not only in probing components of the active nucleus, like the dusty torus that is invoked in unification schemes, but also are they a very useful diagnostic of the (otherwise hard to determine) dust properties and dust composition in other galaxies. Current and future X-ray missions with sensitivity at soft energies and high spectral resolution (like SAX, AXAF and Spektrum-X) will play an important role in studying these issues." + }, + "9802/astro-ph9802238_arXiv.txt": { + "abstract": "This paper summarizes some of our recent projects which try to illuminate the nature and importance of jets associated with active nuclei and compact objects. After a short introduction on jets in radio galaxies and radio loud quasars the paper focuses on jets in other source types, such as radio-quiet quasars, Seyfert and LINER galaxies, and stellar mass black holes. Radio observations of quasars, for example, have brought new evidence for the existence of relativistic jets in radio quiet quasars, while HST and VLA observations of Seyfert galaxies have now clearly established not only the presence of radio jets, but also the great importance these jets have for the morphology and the excitation of the emission line region in these AGN. Moreover, a recent VLA survey found a large fraction of low-luminosity AGN to host compact, flat-spectrum radio cores indicating the presence of radio jets there as well. Finally the jet/disk-symbiosis model, which successfully explains radio cores in LINER galaxies, is applied to the stellar mass black hole GRS~1915+105, indicating that the radio cores in both types of sources are just different sides of the same coin. The conclusion drawn from all these observations is that radio jets are a ubiquitous feature of most---if not all---AGN and play an important effect in the overall energy budget, as well as for the interpretation of observations in other wavebands (e.g. optical emission lines). ", + "introduction": "One of the main subjects for radio astronomers has been the study of extragalactic radio jets. When observed at a higher resolution many of the first radio sources which were discovered in the early years of radio astronomy later turned out to be powerful, collimated plasma flows of relativistic plasma (jets) which were ejected from the nucleus of giant elliptical galaxies. These structures can reach sizes of several million light years and hence extend way beyond their host galaxies into the vastness of intergalactic space. This relative isolation is ideal to study the physics of astrophysical plasma flows in great detail (see e.g.~Bridle \\& Perley 1984, Bridle et al.~1994, Marti et al.~1997) and allows to make some estimates of the properties of the IGM (inter-galactic medium, e.g. Subrahmanyan \\& Saripalli 1993). An even more important aspect of radio jets, however, is that they are the most visible sign, the \"smoking gun\" so to speak, of Active Galactic Nuclei (AGN) which are thought to be powered by accreting black holes. Hence, jets have been studied with great interest over many years and a huge zoo of different jet species has emerged which will be discussed as a useful background for the further discussion in the next section. ", + "conclusions": "Today we can make at least one firm statement, namely that there is no known class of active sources, where one suspects an accreting black hole as the central engine, where jets cannot be found in at least a few members of this class. Consequently {\\it jet formation must be in principle possible under almost all circumstances}, whether there are huge or small luminosities, huge or small Eddington-luminosities, or huge or small sizes, there always seems to be a way to make it possible. Given the many observed stellar (Herbig-Haro) jets, this may even be true if the central object is not a black hole. The question remains, however, whether jets always form in each and every case, i.e. whether each accreting black hole always produces a jet? Even though I suspect the answer to be a strong \"yes\", with only a feeble \"but\", a conclusive answer cannot be given yet, since in many cases our sensitivity is not good enough to securely prove or exclude the existence of jets. Much of the progress in the detection of astrophysical jets in recent years has made use of the superior resolution and sensitivity of the VLA, often pushing its ability to the limits, so that further substantial progress may require a further step in radio technology. For Seyferts, the combination of high resolution radio maps with observations of the narrow emission line regions presented in a number of papers, including this one, has now confirmed earlier claims that outflows are a crucial ingredient to the study of Seyferts and therefore most likely also to AGN physics in general. Still, there is a whole universe in front of us that wants to be explored, as for example the low-luminosity AGN in our neighborhood which are receiving more and more attention in recent years. One result is already clear: jets are here to stay, and their importance should not be underestimated as perhaps done in the past by many astronomers not concerned with radio astronomy. \\bigskip\\noindent {\\it Acknowledgment:} This work was done in collaboration with a number of colleagues who I need to thank for their cooperation, especially A.S. Wilson, C. Simpson, A. Patnaik, W. Sherwood, \\& L.C. Ho, and P.L. Biermann. This summary paper is based in some parts on material reported in Falcke et al.~(1998a\\&b). The research was supported by NASA under grants NAGW-3268, NAGW4700, NAG8-1027, and by the DFG, grant Fa 358/1-1\\&2." + }, + "9802/astro-ph9802242_arXiv.txt": { + "abstract": "OSSE has observed seven transient black hole candidates: GRO~J0422+32, GX339--4, GRS~1716--249, GRS~1009--45, 4U~1543--47, GRO~J1655--40, and GRS~1915+105. Two gamma-ray spectral states are evident and, based on a limited number of contemporaneous X-ray and gamma-ray observations, these states appear to be correlated with X-ray states. The former three objects show hard spectra below 100 keV (photon number indices $\\Gamma < 2$) that are exponentially cut off with folding energy $\\sim$ 100 keV, a spectral form that is consistent with thermal Comptonization. This ``breaking gamma-ray state'' is the high-energy extension of the X-ray low, hard state. In this state, the majority of the luminosity is above the X-ray band, carried by photons of energy $\\sim$100 keV. The latter four objects exhibit a ``power-law gamma-ray state'' with a relatively soft spectral index ($\\Gamma \\sim 2.5-3$) and no evidence for a spectral break. For GRO~J1655--40, the lower limit on the break energy is 690 keV. GRS~1716--249 exhibits both spectral states, with the power-law state having significantly lower gamma-ray luminosity. The power-law gamma-ray state is associated with the presence of a strong ultrasoft X-ray excess (kT $\\sim$ 1 keV), the signature of the X-ray high, soft (or perhaps very high) state. The physical process responsible for the unbroken power law is not well understood, although the spectra are consistent with bulk-motion Comptonization in the convergent accretion flow. ", + "introduction": "The OSSE instrument on the Compton GRO consists of four nearly identical large-area NaI(Tl)--CsI(Na) phoswich detector systems (Johnson et al. 1993). It can be configured to cover the energy range from $\\simeq$40 keV to 10 MeV with good spectral resolution. Simultaneously with the spectroscopy data, count-rate samples in a number of energy bands can be collected on time scales of 4--32 ms to study rapid variability. Hard X-ray transients, including those later classified as BHCs, have been high-priority targets for OSSE throughout the Compton GRO mission. They have generally been observed as targets of opportunity in response to detection of a significant flaring event by the BATSE instrument on the Compton GRO. These target-of-opportunity pointings have lasted from less than 24 hours for GRS~1009--45 to a few weeks for GRO~J0422+32, depending on the strength and duration of the outburst, as well as the flexibility of the Compton observing program. Table \\ref{dates} summarizes the dates on which OSSE observed the seven gamma-ray transients discussed here. Hard X-ray lightcurves (20--100 keV) from BATSE for each transient in this study are shown in Fig. \\ref{lightcurves}. Each data point corresponds to a single day. Note that one of two different flux scales is used in each panel, depending on the intensity of the emission. The total Crab nebular and pulsar emission in this band is $\\simeq$0.32 ph cm$^{-2}$ s$^{-1}$; thus the outbursts studied here peak between about $1/3$ and 1.5 Crab flux units. OSSE observing periods are indicated by shaded vertical intervals in each panel. Note also that the lightcurves end on MJD 49800.0 (1995 Mar 24.0), the last date for which BATSE flux measurements are publicly available as of this writing. For some sources, this is before the last OSSE observation included in this study, and indeed before the first OSSE observation. The lightcurves, nevertheless, serve to place the OSSE measurements in the context of the recent historical behavior of each source. \\subsection{ GRO~J0422+32 } GRO~J0422+32 (XN Per 1992) was discovered in outburst by BATSE in data from 1992 August 5. The lightcurve of the outburst showed a fast rise and approximately exponential decay with $\\tau \\simeq 40$ days, and a secondary maximum beginning $\\simeq$125 days after the onset of the initial outburst (Harmon et al. 1994). It was observed in X-rays at various times in the outburst by ASCA (Tanaka 1993a), ROSAT (Pietsch et al. 1993), and Mir/TTM (Sunyaev et al. 1993). The X-ray spectrum in all cases was consistent with a simple power law, with no evidence for an ultrasoft component. While the mass function of $1.2 \\pm 0.04 \\Msun$ determined by Filippenko, Matheson, \\& Ho (1995) is low enough that the compact object might indeed be a neutron star, the H$\\alpha$ radial velocity curve and the M stellar type of the secondary imply a mass of 3.6$\\Msun$ for the primary. The photometric measurements of Callanan et al. (1996) support this mass estimate and give a distance of $\\sim$2 kpc. It was a transient radio source, exhibiting a radio lightcurve consistent with a synchroton bubble ejection event (Shrader et al. 1994). OSSE observed GRO~J0422+32 for 33 days spanning the interval from the peak of the outburst through the decline to approximately half maximum intensity at 100 keV. There was strong, rapid intensity variability above 50 keV (rms $\\simeq$ 30--40\\% of the total emission) and a substantial peaked noise component in the power spectrum (Grove et al. 1994). The photon spectrum was described well by the simple, exponentially truncated power-law form in Eqn. 1. While this spectral shape was valid for each observation day, there was substantial evolution of the model parameters throughout the outburst. As the intensity at 100 keV declined, the exponential folding energy increased by about 20\\%, with a nearly linear anticorrelation between these parameters. % The TTM and HEXE instruments on Mir-Kvant observed 1992 Aug 29 -- Sep 2 and reported a hard spectrum extending well above 100~keV (Sunyaev et al. 1993; Maisack et al. 1994). To create a consistent broadband spectrum, we therefore selected OSSE data only from a four-day period surrounding this interval and fit the combined spectrum (Table \\ref{fits}). \\subsection{ GRS~1009--45} The transient GRS 1009--45 (XN Vel 1993) was discovered by GRANAT/WATCH on 1993 September 12 (Lapshov, Sazonov, \\& Sunyaev 1993). The lightcurve above 20 keV for the discovery outburst had a fast rise and exponential decay with $\\tau \\simeq 5$ days (Harmon et al. 1994). The secondary outbursts, peaking $\\sim$35 days and $\\sim$85 days after the initial outburst, had substantially longer rise and fall times (see Fig. \\ref{lightcurves}). The X-ray spectrum was ultrasoft (Kaniovsky, Borozdin, \\& Sunyaev 1993; Tanaka 1993b; Moss 1997; Ebisawa 1997). A blue optical counterpart has been identified (Della Valle \\& Benetti 1993), of type late-G or early-K in a $6.86 \\pm 0.12$ h binary orbit (Shahbaz et al. 1996). A mass function has yet to be determined. The distance is estimated to be between 1.5 and 4.5 kpc. OSSE observed GRS 1009--45 for a single 17-hour period $\\simeq$9 days after the onset of the outburst, and $\\simeq$6 days after the peak, by which time the 20--100 keV flux had decayed to $\\sim$1/3 its maximum (Harmon et al. 1994). Despite the brevity of the pointing, OSSE clearly detected emission above 300 keV, with a power law spectrum and no evidence for a break. ASCA data from 1993 November 11 are shown in Fig. \\ref{spectra} together with the OSSE spectrum. The ASCA data, which have been taken from Moss (1997), have been scaled upward by a factor of $\\sim$30, corresponding to the decline in the BATSE hard X-ray flux between the OSSE and ASCA observations. \\subsection{4U~1543--47} BATSE detected 4U~1543--47 on 1992 April 18 in its third known outburst since the discovery in 1971 (Matilsky et al. 1972). The X-ray spectrum in the discovery outburst was ultrasoft. Lightcurves show a fast rise and exponential decay, and multiple secondary maxima have been observed. The e-folding time of the decline of the X-ray emission from the discovery outburst was $\\tau \\simeq$ 85 days, while BATSE measured a much shorter $\\tau \\simeq$ 2.5 days above 20 keV in 1992. A similar difference in soft and hard X-ray decay times was observed from GS~1124--68 (Gilfanov et al. 1991). ROSAT observed 4U~1543--47 in 1992 August and September (Greiner et al. 1993), when the source was undetectable by BATSE. The X-ray spectrum was again ultrasoft. There are no reports in the literature of rapid intensity variations. The optical counterpart is an A-type dwarf at $\\sim$4 kpc in an unknown binary orbit (Chevalier 1989). OSSE observed 4U~1543--47 for a nine-day period beginning about 10 days after its detection by BATSE. The outburst peaked on April 19, and had decayed to undetectability by BATSE by the time the OSSE observation began. The OSSE data indicate that the source underwent a secondary outburst during this observation, at a flux level below BATSE's daily sensitivity. Statistically significant emission was detected by OSSE on the first six observing days, 1992 April 28 -- May 3. The spectrum was a simple power law, with emission detected to at least 200 keV. While the gamma-ray luminosity varied by a factor of $\\sim$3, there was no statistically significant variation in the power-law index from day to day. \\subsection{GRO~J1655--40} The transient X-ray source GRO~J1655--40 (XN Sco 1994) was discovered by BATSE in data from 1994 July 27. Since the discovery outburst, the emission has been episodic, showing periods of weeks of strong, relatively constant emission, then returning to quiescence (Harmon et al. 1995). The mass function is measured to be 3.35$\\pm$0.14 $\\Msun$, and optical eclipses are observed with a 2.6-day period, suggesting that the system is viewed nearly edge-on (Bailyn et al. 1995a,b). Orosz et al. (1996) estimate the mass of the compact object to be 7$\\Msun$. It is a strong, transient radio source with superluminal jets (Tingay et al. 1995, Hjellming \\& Rupen 1995) at a distance of 3.2 kpc (McKay et al. 1994, Hjellming \\& Rupen 1995, Bailyn et al. 1995b). In at least three cases, the X-ray outbursts were followed within days by ejection events in the radio jets. This source has been observed by OSSE on six occasions. Results from the first five were presented by Kroeger et al. (1996). Four observations provided strong detections of a soft power-law spectrum, while in the remaining two observations the source was undetectable by OSSE. The dates of the former are given in Table \\ref{dates}. A simple power-law model gives a good fit to each observation, as well as to the sum of all observations with high gamma-ray luminosity, $L_\\gamma$. Table \\ref{fits} reports the photon index for the sum of the high $L_\\gamma$ observations. The envelope of the scatter of spectral index for the individual observations is 0.4. Investigation of the spectral index on $\\sim$90-minute timescales reveals a weak dependence of index with $L_\\gamma$: higher luminosity intervals tend to have more negative spectral index (i.e. a steeper spectrum). ASCA observed GRO~J1655--40 on 1995 Aug 15-16, during an outburst that unfortunately was not observed by OSSE. The BATSE instrument, with its all-sky capability, monitored this outburst in its entirety. Combined ASCA and BATSE spectra show an ultrasoft excess and a soft power-law tail ($\\Gamma \\simeq 2.4$) extending beyond 100 keV (Zhang et al. 1997). The ASCA data plotted in Fig. \\ref{spectra} have been scaled by the ratio of the average OSSE flux at 40 keV to the reported BATSE flux at 40 keV for 1995 Aug 15-16 (BATSE data are not shown in the figure). The ultrasoft excess remains apparent. During the most recent OSSE observation of GRO~J1655--40 (1996 Aug-Sep), the source was also observed by the Rossi XTE, so high-quality, simultaneous, broadband spectra should soon be available. \\subsection{GX339--4} The highly variable black hole candidate GX339--4 (1H 1659-48.7) exhibits all four of the X-ray states (Miyamoto et al. 1991 and references therein). Its optical counterpart (Doxsey et al. 1979) lies at a distance generally assumed to be $\\sim$4 kpc. It is a variable radio source (Sood \\& Campbell-Wilson 1994), and recent evidence suggests it may have a jet (Fender et al. 1997). In the X-ray low, hard state, rapid fluctuations with rms variations as high as 30\\% have been observed, as well as QPOs near $\\sim$0.8 Hz (Grebenev et al. 1991), $\\sim$0.1 Hz and $\\sim$0.05 Hz (Motch et al. 1983). Observations by Ginga in 1991 September established that the source was in the X-ray low, hard state. Contemporaneous OSSE observations revealed that the low-energy gamma-ray emission was strong and that the spectrum was exponentially cut off (see Fig. \\ref{spectra} and Table \\ref{fits}), from which Grabelsky et al. (1995) argued that the X-ray low, hard state corresponds to gamma-ray outburst in this object. Observations performed with GRANAT (Grebenev et al. 1993) indeed show that GX339--4 exhibits distinct hard X-ray states correlated with the X-ray behavior. OSSE has observed GX339--4 four times. The first pointing, in 1991 September, was made as a target of opportunity in response to an outburst detected by BATSE. The second observation was performed in 1991 November as part of the scheduled viewing plan, when the source happened to have very low gamma-ray luminosity (Fig. \\ref{lightcurves}). Because of the potential for confusion with diffuse continuum emission from the galactic plane, we have not included these data in our analysis. The final two observations contain proprietary data and have not been included in our analysis or in Table \\ref{dates}. \\subsection{GRS~1716--249} The transient GRS 1716--249 (GRO~J1719--24, XN Oph 1993) was discovered by GRANAT/SIGMA on 1993 September 25 (Ballet et al. 1993). Observations by ASCA (Tanaka 1993c; Moss 1997; Ebisawa 1997) and Mir-Kvant (Borozdin, Arefiev, \\& Sunyaev 1993) showed a hard power law spectrum, with no evidence for an ultrasoft excess. Rapid intensity variations were detected by BATSE (Harmon et al. 1993) and ASCA. A strong QPO peak was apparent, with a centroid frequency that varied from $\\sim$0.04 Hz to $\\sim$0.3 Hz through the discovery outburst (van der Hooft et al. 1996). A radio (Mirabel, Rodriguez, \\& Cordier 1993) and optical (Della Valle, Mirabel, \\& Cordier 1993) low-mass counterpart has been identified at an estimated distance of $\\simeq$2.4 kpc (Della Valle, Mirabel, \\& Rodriguez 1994). Interpreting optical modulations at 14.7 h as a superhump period, Masetti et al. (1996) estimated the mass of the compact object to be $>$4.9$\\Msun$. Hjellming et al. (1996) reported a rapidly decaying radio flare following the final outburst shown in Fig. \\ref{lightcurves}. This temporal coincidence suggests a connection between the hard X-ray decline and the ejection of the relativistic electrons responsible for the radio emission, and from the similarity of this event to the emergence of a relativistic radio-synchrotron jet in GRO~J1655--40 (Hjellming \\& Rupen 1995), Hjellming et al. suggested that a relativistic jet might have been present in GRS~1716--249 at this time. The radio data were inadequate to confirm or reject this suggestion, unfortunately. OSSE observed GRS 1716--249 five times, with the first two observations beginning about 30 days after the onset of the discovery outburst, while the 20--100 keV flux was still within 10--15\\% of its maximum intensity (these observations are closely spaced in time and indistinguishable in Fig. \\ref{lightcurves}). The remaining three observations were made on the trailing edge of the second outburst, between the second and third outbursts, and finally near the peak of the third outburst. The gamma-ray luminosity varied by an order of magnitude between observations. There is good evidence that both spectral states were observed. The first, second, and fifth observations, all made near the peak of outbursts, show the exponentially truncated power law form of Eqn. 1, while the third is a simple power law spectrum. The spectrum of the fourth observation can be fit with either functional form, and a statistically significant distinction cannot be drawn. The exponentially truncated spectra have been summed and are shown in the upper GRS~1716--249 spectrum in Fig. \\ref{spectra}, while the third observation is shown in the lower spectrum. The spectral parameters in Table \\ref{fits} are similarly divided. OSSE has detected spectral shape changes only in this BHC and Cyg X-1 (Phlips et al. 1996). Also plotted in Fig. \\ref{spectra} is the X-ray spectrum from an ASCA observation on 1993 October 5 (Moss 1997), approximately three weeks before the first OSSE observation and during the discovery outburst. No ultrasoft excess is apparent, which indicates that the source is in the X-ray low, hard state. \\subsection{GRS~1915+105} The transient GRS 1915+105 (XN Aql 1992) was discovered by GRANAT/WATCH on 1992 August 15 (Castro-Tirado et al. 1992). The discovery outburst began with a slow rise to maximum intensity, and since that time its hard X-ray lightcurve has been episodic (Fig. \\ref{lightcurves}). On timescales of tens to thousands of seconds, the X-ray emission is dramatically variable, showing several repeating temporal structures (Greiner, Morgan, \\& Remillard 1996). The X-ray spectrum is also variable: during intense flares the spectrum is cut off near 5 keV, while during more stable periods it shows an ultrasoft disk blackbody plus a power law with $\\Gamma = -2.2$ (Greiner et al. 1996). At a distance of $\\simeq$12.5 kpc, it was the first object in our galaxy shown to have superluminal radio jets (Mirabel \\& Rodriguez 1994). As with GRO~J1655--40, radio flares frequently, although not always, follow hard X-ray flares and occur during a decrease in the hard X-ray emission (Foster et al. 1996). Mirabel et al. (1997) have suggested that the companion is a high-mass late Oe or early Be star. A mass function for the system is not yet available. OSSE observed GRS~1915+105 on three occasions, all of which are after the final day shown in Fig. \\ref{lightcurves}. Extrapolating the best-fit power law model parameters from each of these viewing periods into the 20--100 keV band gives 0.09, 0.07, and 0.07 ph cm$^{-2}$ s$^{-1}$, respectively, indicating that both observations were near the historical peak hard X-ray intensity for this source (compare Fig. \\ref{lightcurves}). In all cases, the spectrum is a simple power law with no evidence for a break. As with GRO~J1655--40, there is a weak anti-correlation of spectral hardness with $L_{\\gamma}$ (i.e. as $L_{\\gamma}$ increases, the spectrum becomes softer and the photon number index becomes more negative). GRS~1915+105 was observed in 1996 October by ASCA, Rossi XTE, Beppo SAX, and OSSE in a scheduled campaign during a continuing outburst, so simultaneous, high-quality, broadband spectra are forthcoming for this object. \\subsection{Spectral Analysis} Photon number spectra from each BHC are shown in Fig. \\ref{spectra}, along with the best-fit analytic model extrapolated to 10 keV. In most cases, the spectra are averaged over the entire OSSE dataset for which the source was detected. GRS~1716--249 appears in this figure twice, with the high-luminosity breaking spectra summed together, and the low-luminosity power-law spectra summed together, as discussed below. For clarity of the figure, the spectra have been multiplied by the arbitrary scaling factors indicated next to the source name. X-ray data are shown for GRO~J0422+32, GRS~1009--45, GRO~J1655--40, and GRS~1716--249. We fit the average spectra using one of two general analytic spectral models, either a simple power law, or a power law that is exponentially truncated above a break energy: \\begin{equation} f(E) = \\left\\{ \\begin{array}{ll} A \\, E^{-\\Gamma} \t\t\t\t& \\mbox{$E < E_b$} \\\\ A \\, E^{-\\Gamma} \\, \\exp(-(E-E_b)/E_f) \t& \\mbox{$E > E_b$} \\end{array} \\right. \\end{equation} where $A$ is the photon number flux, $\\Gamma$ is the photon number index, $E_b$ is the break energy, and $E_f$ is the exponential folding energy. Below the break energy, the exponential factor is replaced by unity, and the model simplifies to a power law. Best-fit model parameters are given in Table \\ref{fits}, along with the corresponding luminosities in the gamma-ray band (i.e. above 50 keV). Uncertainties in the model parameters are statistical only and reported as 68\\%-confidence intervals or 95\\%-confidence lower limits. If the simple power law is a statistically adequate fit, we report in Table \\ref{fits} the photon number index from that model along with lower limits to $E_b$ and $E_f$. Because of the strong correlation between these two parameters when neither is required by the data, to establish the lower limits we fixed $E_f = 2 E_b$, since that relation roughly holds for GRO~J0422+32. For four sources --- GRS~1009--45, 4U~1543--47, GX339--4, and GRS~1915+105 --- little change in spectral parameters is observed from day to day or from observation to observation when the source is detected. Larger changes are observed for GRO~J0422+32 (where $E_f$ varies), GRO~J1655--40 (where $\\Gamma$ varies), and GRS~1716--249 (where the spectral shape changes). In Table \\ref{fits} we also report the range of statistically significant variability from day to day in one spectral parameter, $E_f$ for the breaking spectra and $\\Gamma$ for the power-law spectra, along with the range of luminosities above 50 keV observed from day to day. Two general spectral states in the gamma-ray band are apparent in Fig. \\ref{spectra}. Three transients (GRO~J0422+32, GX339--4, and GRS~1716--249) have breaking spectra. The X-ray photon number index is $\\sim$1.5, and the exponential folding energy is $\\sim$100 keV. The remaining four transients (GRS~1009--45, 4U~1543--47, GRO~J1655--40, and GRS~1915+105) show pure power law spectra, with a softer photon number index, in the range $\\simeq$2.5--3. At low luminosity, GRS~1716-249 also appears to show a pure power law spectrum, with $\\Gamma \\simeq 2.5$. The two spectral states have their peak luminosities in distinct energy ranges. Fig. \\ref{nufnu} shows combined X-ray and OSSE gamma-ray spectra for several sources, plotted as luminosity per logarithmic energy interval. The sources shown in the breaking gamma-ray state, GRO~J0422+32 and GRS~1716--249, have the bulk of their luminosity emitted near 100 keV, while for GRO~J1655--40 in the power-law gamma-ray state, the ultrasoft component is dominant. We also note that the distinction between the spectral states is more apparent plotted in this manner. We have searched for evidence of line emission at the energies suggested in previous observations of BHCs, i.e. at 170 keV from backscattered 511 keV radiation, near 480 keV as reported from SIGMA data on Nova Muscae (Goldwurm et al. 1992, Sunyaev et al. 1992), and at 511 keV. With OSSE's very high sensitivity (roughly one order of magnitude greater than SIGMA's at 511 keV), previously reported lines would be detected at high significance. For example, a broadened 480 keV line at $6 \\times 10^{-3}$ ph cm$^{-2}$ s$^{-1}$, the intensity reported from Nova Muscae, would have been detected by OSSE at $\\sim$40$\\sigma$ in an average 24-hour period. As the SIGMA data suggest, these lines might indeed be variable, so we have searched our data on several timescales. Table \\ref{lines} summarizes the results of our search, reporting 5$\\sigma$ upper limits daily, weekly, and on the total of all observations. (We chose this high statistical confidence level because of the large number of intervals searched.) The observations in Table \\ref{dates} were divided into weeks of 7 contiguous days if possible; otherwise, if the number of contiguous observing days was less than 7, those contiguous days were defined as one ``week.'' No statistically significant narrow or moderately broadened line emission is observed at any time from any of these objects. ", + "conclusions": "We have demonstrated the existence of two distinct spectral states of gamma-ray emission from galactic BHCs. The first state exhibits a relatively hard power-law component with photon index roughly 1.5 that breaks at $\\simeq$50 keV and is exponentially truncated with an e-folding energy of $\\simeq$100 keV. This ``breaking gamma-ray'' is firmly identified with the X-ray low, hard state, i.e. with the absence or weakness of an ultrasoft X-ray component. Sources exhibiting this spectral signature include GRO~J0422+32 and GRS~1716--249, as well as GX~339--4 and Cyg X-1 in their X-ray low, hard states. The second state exhibits a relatively soft power-law gamma-ray spectrum with photon index roughly 2--3 and no evidence of a break. In the case of GRO~J1655--40, the lower limit on the break energy is 690 keV, above the electron rest mass. This ``power law gamma-ray'' state is identified with the presence of an ultrasoft X-ray component, which is the signature of the X-ray high, soft state and the X-ray very high state. Because of the paucity of simultaneous X-ray and gamma-ray observations and the difficulty of distinguishing between the two X-ray high states, the identification of the power law gamma-ray state with which of the two high states is somewhat uncertain, although Cyg~X-1 observations indicate that the correspondence is with the high, soft state (Cui et al. 1997). Sources exhibiting the power law gamma-ray state include GRS~1009-45, 4U~1543--47, GRO~J1655--40, and GRS~1915+105, as well as Cyg~X-1 in its X-ray high, soft state. Gamma-ray spectral state changes have now been observed in GRS~1716--249 and Cyg~X-1. Further simultaneous X-ray and gamma-ray spectroscopy measurements of galactic BHCs will help elucidate the details of the relationship of the emission in the two bands." + }, + "9802/astro-ph9802304_arXiv.txt": { + "abstract": "We report the discovery with BeppoSAX-WFC of two new X-ray sources that were only seen during bursts: SAX~J1753.5-2349 and SAX~J1806.5-2215. For both sources, no steady emission was detected above an upper limit of 5~mCrab (2 to 8 keV) for 3~10$^5$~s around the burst events. The single burst detected from SAX~J1753.5-2349 shows spectral softening and a black body color temperature of 2.0~keV. Following the analogy with bursts in other sources the burst very likely originates in a thermonuclear flash on a neutron star. The first of two burst detected from SAX~J1806.5-2215 does not show spectral softening and cannot be confirmed as a thermonuclear flash. ", + "introduction": "The Wide Field Camera instrument (WFC, Jager et al. 1997) on board the BeppoSAX satellite (e.g., Boella et al. 1997) has the largest field of view (FOV) of any astrophysical X-ray imaging device flown so far. This implies an exceptional capability in finding short unexpected transient sources with durations between seconds and one hour, such as gamma and X-ray bursts. Only instruments that cover relevant portions of the sky with an appropriate time coverage can be efficient in finding such events. The relevant sky portion for gamma-ray bursts is the whole sky, for X-ray bursts this is the galactic bulge. According to a recent count of X-ray bursters before the launch of BeppoSAX (Van Paradijs 1995), about 80\\% of all known X-ray bursters are within 20 degrees from the direction to the galactic center. WFC can cover all of this region with a single pointing and it is obvious it can contribute a lot to the knowledge of sources of X-ray bursts. The galactic bulge region is regularly monitored with WFC as part of an ongoing program. So far, there have been campaigns in the spring of 1997 and the falls of 1996 and 1997. In the first operational year of BeppoSAX over 10$^6$~s has been accumulated on the region and has turned up a number of new X-ray bursters (In~'t~Zand et al., 1998a, In~'t~Zand et al. 1998b, Heise et al. 1997, Cocchi et al. 1997a). Four of these were first detected as X-ray sources. This paper presents two of those. They are set apart by the fact that they were only detected during bursts, no steady emission was detected. A general review of the bulge observations with WFC may be found in Heise~et~al.~1997. The WFC instrument consists of two identical coded aperture cameras. Each has a FOV of 40 by 40 square degrees and covers 3.7\\% of the sky. The angular resolution is 5' full width at half maximum. The cameras are pointed in opposite directions. Because of the low-earth orbit of BeppoSAX the FOV of either camera is at any time usually blocked by the earth. The other X-ray instruments, the so called narrow field instruments, view the sky perpendicular to the WFC. Most of the galactic bulge observations with WFC are performed with WFC as prime instrument and the center of the FOV of either camera is pointed as close as possible to the galactic center. The very bright and, therefore, disturbing source Sco~X-1 is usually kept outside the FOV. \\begin{figure}[t] \\begin{center} \\leavevmode \\epsfxsize=7cm \\epsfbox{fig1753lc.ps} \\caption{ Time history of SAX~J1753.5-2349 in two bandpasses and in the total bandpass. The photons have been counted in the portion of the detector illuminated by the source. No background subtraction was performed. The bin time is 1~s.\\label{fig1753lc} } \\end{center} \\end{figure} \\begin{figure}[t] \\begin{center} \\leavevmode \\epsfxsize=7cm \\epsfbox{fig1806lc.ps} \\caption{ Time history of SAX~J1806.5-2213 in two bandpasses and in the total bandpass. See caption to Figure~\\ref{fig1753lc} for explanation. The bin time is 3~s.\\label{fig1806lc} } \\end{center} \\end{figure} \\begin{figure}[t] \\begin{center} \\leavevmode \\epsfxsize=7cm \\epsfbox{fig1753err.ps} \\caption{ Error region of SAX~J1753.5-2349 for a confidence level of 99\\%. The cross indicates the best fit position \\label{fig1753err} } \\end{center} \\end{figure} \\begin{figure}[t] \\begin{center} \\leavevmode \\epsfxsize=7cm \\epsfbox{fig1806err.ps} \\caption{ Error region of SAX~J1806.5-2215 for a confidence level of 99\\%, based on the first burst. The cross indicates the best fit position \\label{fig1806err} } \\end{center} \\end{figure} ", + "conclusions": "" + }, + "9802/astro-ph9802169_arXiv.txt": { + "abstract": "We have created a map of the large-scale infrared surface brightness in excess of that associated with the atomic interstellar medium, using region-by-region correlations between the far-infrared and 21-cm line surface brightness. Our study updates and extends a previous attempt with the {\\it Infrared Astronomical Satellite} and Berkeley/Parkes H~I surveys. The far-infrared observations used here are from the {\\it Cosmic Background Explorer} Diffuse Infrared Background Experiment, which extends far-infrared wavelength coverage to 240 $\\mu$m, so that we are reliably sampling the emission of large, thermal-equilibrium grains that dominate the dust mass. The H~I data are from the combined Leiden-Dwingeloo and Parkes 21-cm line surveys. Using the maps of excess infrared emission at 100, 140, and 240 $\\mu$m, we created an atlas and identified the coherent structures. These infrared excess clouds can be caused both by dust that is warmer than average, or by dust associated with gas other than the atomic interstellar medium. We find very few warm clouds, such as the H~II region around the high-latitude B-type star $\\alpha$ Vir. The majority of the infrared excess clouds are {\\it colder} than the average atomic interstellar medium. These clouds are peaks of column density, and their excess infrared emission is due to dust associated with molecular gas. We identify essentially all known high-latitude molecular clouds in the infrared excess maps, and further identify a sample of new clouds with similar infrared properties. The infrared excess was correlated with CO line brightness, allowing us to measure the ratio of H$_2$ column density to CO line integral (i.e. the $N$(H$_2$)/$W$(CO) conversion factor) for high-latitude clouds. The atlas of infrared excess clouds may be useful as a guide to regions of relatively higher interstellar column density, which might cause high extinction toward extragalactic objects at optical and ultraviolet wavelengths and confusion toward structures in the cosmic background at infrared and microwave frequencies. ", + "introduction": "It has been known for some time that interstellar molecular gas exists away from the galactic midplane, based on the presence of absorption lines of H$_2$ and other molecules in the spectra of high-latitude stars (\\cite{savage77}) and the presence of millimeter-wave emission of CO from regions of high optical extinction (Blitz, Magnani, \\& Mundy 1984). However, the pervasiveness of the molecular component of the local interstellar medium has not been convincingly assessed, because of the difficulties of observing molecular gas over large areas. Absorption line observations are restricted either to trace elements or to lines of sight toward the few hot, bright stars at high latitude, because the electronic ground-state transitions of both atomic and molecular hydrogen are in the ultraviolet. Millimeter-wave observations of CO are restricted to relatively small areas because all-sky observations are at present prohibitively expensive. An early survey of randomly-selected lines of sight found a surface filling fraction of 0.5\\% for CO($1\\rightarrow 0$) emission (Magnani, Lada, \\& Blitz 1986). A large---but, necessarily, incompletely-sampled---survey was recently performed with a very low detection rate of 0.3\\%, suggesting that the northern galactic hemisphere is largely `devoid' of molecular gas (Hartmann, Magnani, \\& Thaddeus 1997). On the other hand, there remains the question of whether a survey in a particular observable quantity, such as the brightness of collisionally-excited rotational line emission of a particular molecule, is really tracing all of the molecular gas. The H$_2$ and CO might have somewhat different spatial distributions due to their different photodissociation cross-sections and chemistry, and the rotational energy levels of CO may not be excited in low-density ($n({\\rm H}_2)<10^3$ cm$^{-3}$) environments. With the advent of the infrared all-sky survey by the {\\it Infrared Astronomical Satellite} (\\IRAS), a new light was cast on study of the interstellar medium (\\cite{neugebauer,low}). The 100 $\\mu$m surface brightness was found to be well-correlated with the H~I column density on both small and large scales (\\cite{bp88}), demonstrating its value as a tracer of interstellar gas. If the infrared emission arises from both the atomic and molecular phases of the interstellar medium, then regions with an excess infrared emission relative to the H~I column density are likely locations of molecular gas. This effect has been demonstrated in Ursa Major (\\cite{vht87}) and for isolated cirrus clouds (Heiles, Reach, \\& Koo 1988; Reach, Koo, \\& Heiles 1994); in both cases the infrared excess was found to be associated with CO emission. Because the infrared surface brightness and H~I column density are known over the entire sky, it is possible to use the difference between the infrared map and an appropriately-scaled map of H~I column density to produce an all-sky survey of molecular gas. This idea has been exploited by D\\'esert et al. (1988), who used the \\IRAS\\ 100 $\\mu$m data and the Berkeley H~I survey (\\cite{hh}) to create a catalog of infrared excess clouds. In the present paper, we study the distribution and nature of infrared excess clouds using relatively recent data from the {\\it Cosmic Background Explorer}\\footnote{ The National Aeronautics and Space Administration/ Goddard Space Flight Center (NASA/GSFC) is responsible for the design, development, and operation of the Cosmic Background Explorer (\\COBE). Scientific guidance is provided by the \\COBE\\ Science Working Group. GSFC is also responsible for the development of the analysis software and for the production of the mission data sets.} (\\COBE) mission (\\cite{boggess92}) and the Leiden-Dwingeloo H~I survey (\\cite{hartmann97}). These new surveys provide higher sensitivity and higher reliability than the previous infrared and H~I observations for large-scale emission. More importantly, the \\COBE\\ observations at 100, 140 and 240 $\\mu$m wavelength are a reliable measure of the emission from large, thermal-equilibrium grains that dominate the dust mass. From a detailed study of the infrared emission in the Orion region, it was shown that the 100, 140, and 240 $\\mu$m emission sample essentially the same dust temperature along the line of sight (\\cite{Wall}). Our first results (\\cite{reachbaas}), based on comparing the \\COBE\\ 240 $\\mu$m optical depth to the Berkeley H~I surveys, encouraged us to pursue a more thorough comparison of the H~I and infrared data. The large-scale distribution of molecular gas is important for a number of practical applications. \\noindent {\\it Extinction of extragalactic objects---} Extragalactic observations at visible and shorter wavelengths are affected by extinction even at high galactic latitude. In order to estimate the extinction, which affects both the brightness and color of extragalactic objects, it has been necessary to rely on 21-cm line surveys (cf. \\cite{burstein}). If an extragalactic object lies behind a high-latitude molecular cloud, its extinction will be significantly underestimated using the 21-cm line surveys. The results derived here will help observers to identify regions where anomalously high extinction from molecular clouds can be expected. A recent effort by another group (Schlegel, Finkbeiner, \\& Davis 1997) systematically addresses this issue, using the \\COBE\\ and \\IRAS\\ data to create a map of the extinction. Large-scale variations in the temperature and gas-to-dust ratio were calibrated by Schlegel et al. in a manner very similar to ours. By calculating the column density using large-scale average dust temperature, it is likely that the Schlegel et al. dust maps will underestimate the extinction toward relatively cold clouds such as the molecular clouds we have identified. We therefore recommend that the Schlegel et al. maps be used to estimate the extinction, and that our molecular cloud atlas be used as a supplement, warning of cold, high-extinction clouds. Lines of sight behind these clouds should certainly be avoided unless 1-3 magnitudes ($A_V$) of extinction can be tolerated. \\noindent {\\it Shadowing of distant X-rays---} Interstellar clouds produce distinct shadows on the soft X-ray emission from the Galactic halo and the extragalactic emission (\\cite{mccammon}). These shadows have been compared to 21-cm line maps and to \\IRAS\\ 100 $\\mu$m surface brightness maps (Snowden, McCammon, \\& Verter 1993; \\cite{wang95}). In both cases, the interstellar column density is underestimated in the presence of molecular clouds, because molecular gas can absorb X-rays and molecular clouds are relatively faint at 100 $\\mu$m (and nearly invisible at 60 $\\mu$m; \\cite{Laureijs96}), compared to atomic gas. We provide both a map of infrared excess clouds and a calibration of their column density in the present paper. \\noindent {\\it Diffuse $\\gamma$-ray emission and the N(H$_2)/W($CO$)$ factor---} A significant source of the $\\gamma$-ray surface brightness of the sky is due to interaction of cosmic rays with interstellar gas. The expected near-linear proportionality of $\\gamma$-rays with total column density allows a calibration of the molecular column density, when correlated with H~I and CO maps (\\cite{strong88}; \\cite{digel96}). In the presence of H$_2$ clouds with relatively little CO emission, this calibration can be significantly biased. Here we calibrate the molecular column density using the infrared excess, and compare the results to $\\gamma$-ray studies. The ratio of H$_2$ column density to CO line integral is important for assessing the vertical mass distribution of molecular gas in the disk of our Galaxy. \\noindent {\\it Relation to external galaxies---} Low-metallicity galaxies such as the Large and Small Magellanic clouds contain relatively fewer giant molecular clouds, and more translucent regions---due to the lack of dust. High-latitude molecular clouds such as studied here share empirical similarities with low-metallicity galaxies, and should provide insight into the nature of the latter. ", + "conclusions": "The infrared excess clouds, which were found after removing emission associated with the atomic interstellar medium, are predominantly molecular clouds. The dust in these clouds is colder than dust in diffuse atomic clouds, for the same reason that the gas is molecular: the outer layers of the cloud shield the center from the dust-heating and molecule-dissociating effects of the interstellar radiation field. These molecular clouds fill only a small fraction of the high-latitude sky. The column density of the infrared excess clouds can be estimated using the calibrations discussed above, which may be rewritten as \\begin{equation} I_{100}^{ex}/N({\\rm H}_2) = (0.26\\pm0.05) \\times {\\rm 10}^{-20}\\,{\\rm cm}^{2} {\\rm MJy}\\,{\\rm sr}^{-1}. \\end{equation} We can estimate the mass surface density of the infrared excess clouds, using the difference between positive and negative excess for $|b|>20^\\circ$ (Fig.~\\ref{fig:exhis}). The surface density is about 0.3~$M_\\odot$~pc$^{-2}$ (including 40\\% He by mass), roughly equally partitioned among regions of low and high brightness. (We adjusted for the part of the sky at $|b|<20^\\circ$ by assuming it has the same average infrared excess as the rest of the sky.) The mass density that we find is an order of magnitude larger than that found in an unbiased, wide-field CO survey (\\cite{hartmag}) of the northern sky, but it is comparable to a previous estimate---also based on CO---of the total mass of the known molecular clouds at high galactic latitude (\\cite{magnani97}). We suspect that the unbiased survey missed much of the molecular gas because either there was no CO mixed with the diffuse H$_2$ or the CO rotational levels were subthermally excited (because of low H$_2$ volume density). We do {\\it not} find a great disparity between the northern and southern galactic hemispheres (cf. \\cite{hartmag}); according to the infrared excess, the southern galactic hemisphere has about 30\\% more molecular gas than the northern hemisphere. The high-latitude infrared-excess clouds comprise a significant fraction of the total mass of molecular gas in the solar neighborhood. The infrared excess maps can be used to provide a guide to regions of anomalously high extinction. The extinction calculated from the H~I column density (cf. \\cite{burstein}) can significantly underestimate the total extinction in these regions. If we assume that the dust in the infrared-excess clouds has the same extinction cross-section per H-nucleus as diffuse clouds for which this quantity has been measured (Bohlin, Savage, \\& Drake 1978), then the extinction due to infrared excess clouds can be obtained from Fig.~\\ref{fig:exmaps}, after multiplying the 100 $\\mu$m surface brightness by $A_V/I_{100}^{ex}\\sim 0.027$~mag~MJy$^{-1}$~sr." + }, + "9802/astro-ph9802219_arXiv.txt": { + "abstract": "We explore the relationship between the metallicity of the intracluster gas in clusters of galaxies, determined by X-ray spectroscopy, and the presence of cooling flows. Using ASCA spectra and ROSAT images, we demonstrate a clear segregation between the metallicities of clusters with and without cooling flows. On average, cooling-flow clusters have an emission-weighted metallicity a factor $\\sim 1.8$ times higher than that of non-cooling flow systems. We suggest this to be due to the presence of metallicity gradients in the cooling flow clusters, coupled with the sharply peaked X-ray surface brightness profiles of these systems. Non-cooling flow clusters have much flatter X-ray surface brightness distributions and are thought to have undergone recent merger events which may have mixed the central high-metallicity gas with the surrounding less metal-rich material. We find no evidence for evolution in the emission-weighted metallicities of clusters within $z \\sim 0.3$. ", + "introduction": "X-ray spectroscopy provides an accurate measure of the metallicity of the hot intracluster medium (ICM) in clusters of galaxies. The strengths of the various emission lines relative to the continuum reveal the abundances of the emitting elements relative to hydrogen. Such measurements are particularly clear for the case of iron in rich clusters where the K-shell lines are typically well-defined in the X-ray spectra. The mass of the ICM dominates over the visible mass in stars in a cluster by a factor of 2--5 (Arnaud \\etal 1992; David, Jones \\& Forman 1995; White \\& Fabian 1995). Metallicity measurements from X-ray observations thus provide firm constraints on the history of metal production within the cluster potential wells. For low-redshift clusters, the observed abundance of iron in the ICM is approximately 1/3 solar (\\eg Edge \\& Stewart 1991). The strong correlation between the mass of iron and the total optical light from elliptical and lenticular galaxies within a cluster suggests that early-type galaxies are responsible for the bulk of the enrichment (Arnaud 1992). Mushotzky \\etal (1996) showed that the relative abundances of individual elements such as Si, S and Fe determined from X-ray spectra suggest that most of the metals originate from type II supernovae. Mushotzky and Lowenstein (1997) further demonstrated that the iron abundance in rich clusters shows little evolution between $z \\sim 0.3$ and now, suggesting that most of the enrichment of the ICM occurred at high redshifts ($z > 0.3$). This is consistent with current semi-analytic models of galaxy formation (e.g. Kauffmann \\& Charlot 1997) which find that more than 80 per cent of the metal enrichment occurs at $z > 1$. X-ray observations of clusters of galaxies show that in the central regions of most ($70-90$ per cent) clusters the cooling time of the ICM is significantly less than the Hubble time (Edge \\etal 1992; Peres \\etal 1997). The observed cooling leads to a slow net inflow of material towards the cluster centre; a process known as a cooling flow (Fabian 1994). The X-ray imaging data show that gas typically `cools out' throughout the central few tens to hundreds of kpc in the clusters. Recent spatially resolved X-ray spectroscopy has confirmed the presence of distributed cool (and rapidly cooling) gas in cooling flows, with a spatial distribution and luminosity in excellent agreement with the predictions from the imaging data (Allen \\& Fabian 1997). The cooling flow in an X-ray luminous cluster can account for up to $\\sim 70$ per cent of the total bolometric luminosity of the system (about half of this luminosity being due to material cooling out of the flow, and the rest due to the gravitational work done on the gas as it flows inwards; Allen \\etal 1998). Where abundance measurements have been compiled for large samples of clusters, a significant dispersion in the metallicity (which is primarily determined from the iron abundance) for clusters of a fixed X-ray luminosity and/or temperature has been revealed. In particular, the metallicity appears to depend on whether or not a cluster has a cooling flow (Yamashita 1992; Fabian \\etal 1994a) in the sense that cooling-flow clusters tend to have higher metallicities. Fabian \\etal (1994a) speculated that this may be due to differences in the origins of the clusters, to the intracluster gas in the cooling-flow clusters being inhomogeneous with the cooler gas being more metal-rich, or due to the presence of metallicity gradients in these systems. Those clusters without cooling flows are assumed to have experienced a major merger event which has mixed the gas, leading to a lower {\\it emission-weighted} metallicity. In this paper we present metallicity measurements for a sample of 30 X-ray luminous clusters observed with ASCA and ROSAT (see Allen \\etal 1998 for details). We identify the cooling flows in the sample and explicitly account for the effects of the cooling gas on the X-ray spectra. We argue that an abundance gradient in the cooling flow clusters underlies the variation seen in the emission-weighted metallicities. Throughout this paper, we assume $H_0$=50 \\kmpspMpc, $\\Omega = 1$ and $\\Lambda = 0$. ", + "conclusions": "\\subsection{Metallicity gradients, mergers and mixing} Fabian \\etal (1994a) suggested that the higher emission-weighted metallicities for CF clusters could be due to inhomogeneities in the ICM, with small blobs of cooler, denser and more metal-rich gas being immersed in hotter gas at the centres of the CF systems. The distributed mass deposition profiles within cooling flows requires that the ICM there is inhomogeneous (Fabian 1994 and references therein). Fabian \\etal (1994a) also noted the potential importance of abundance gradients in clusters. Abundance gradients appear to be a common feature of the cores of CF clusters. First noticed in ASCA (Fukazawa \\etal 1994) and ROSAT PSPC (Allen \\& Fabian 1994) spectra of the Centaurus cluster, an increase in the iron abundance towards the centres of CF clusters has since been found in ASCA spectra for Abell 496 (Hatsukade 1997), the Virgo cluster (Matsumoto \\etal 1996) and AWM7 (Ezawa \\etal 1997). The abundance of iron across the Perseus cluster, which also hosts a large cooling flow, appears patchy with the highest value probably occurring in the core (Arnaud \\etal 1994; see also Molendi \\etal 1998). In contrast, ASCA studies of Abell 1060 (Tamura \\etal 1996), the Ophiuchus cluster (Matsuzawa \\etal 1996) and the Coma Cluster (Watanabe \\etal 1997), which have little or no cooling flows, show no abundance gradients. An abundance gradient will strongly influence the emission-weighted metallicity determined from the integrated X-ray spectrum of a cluster. The X-ray emission depends on the square of the ICM density and so is dominated by the innermost, densest regions of a cluster. This is particularly the case for CF clusters, where a significant fraction (up to $\\sim 70$ per cent) of the total X-ray luminosity may arise from within the central $2-3$ hundred kpc (Peres \\etal 1997; Allen \\etal 1998). We have simulated the effects of metallicity gradients on the mean emission-weighted metallicities for CF and NCF clusters using the observed X-ray surface brightness profiles for Abell 478 and 2218 as representative examples of CF and NCF systems. (These clusters have the highest-quality imaging data.) We find that a linear gradient in metallicity dropping from $\\sim 0.8$ solar at the centre to $\\sim 0.1$ solar at 500 kpc (and remaining roughly constant outside this radius) leads to a mean emission-weighted metallicity, for a CF cluster like Abell 478, of about 0.4. For a NCF system, for which the X-ray surface brightness profile will be much less sharply-peaked, the mean emission-weighted metallicity will be only $\\sim 0.2$. Such gradients should be easily detectable with AXAF. The abundance gradient model provides the most natural explanation for the segregation in the metallicity results for the CF and NCF clusters. The metal-rich core may be the oldest part of a cluster, forming earliest in the deep potential wells. This could enhance both the formation of massive stars and the retention of gas in these regions. (We note that many of the most-massive CF clusters also show evidence for ongoing star formation in their cores; \\eg Allen 1995). After their initial collapse, clusters continue to evolve by the accretion of material, often via subcluster merger events. NCF clusters, such as those included in this study, are thought to have recently experienced a major merger event wherein a large mass component has strongly interacted with the cluster core (Allen 1998). Such events will significantly disrupt the X-ray gas in the core regions of the clusters and will mix and spread the central high-metallicity gas with the outer less metal-rich material. The abundance gradients in NCF clusters are thereby reduced (or even destroyed), although the {\\it total} mass of metals in the ICM is unchanged. (An implication of this is that the global metallicity of the ICM is more accurately estimated for the NCF systems, where the metals are more evenly mixed with the cluster gas.) The dependence of cluster metallicities on the presence or absence of a cooling flows then reflects whether these systems have had their core regions left undisturbed, or recently mixed, rather than on any internal property of the cooling flows such as their mass deposition rates. Reisenegger \\etal (1996) have suggested that the abundance gradient in the Centaurus cluster could result from the cooling flow in that cluster concentrating the metals ejected by type-Ia supernovae in the outer regions of the cD galaxy. Although this mechanism is plausible for low-luminosity systems like the Centaurus cluster, with relatively high ratios of the stellar mass in the central galaxy to the X-ray gas mass, it has more difficulty in accounting for the gradients in more X-ray luminous clusters, with lower stellar/X-ray gas mass ratios. This implies that the metallicity gradients in luminous CF clusters may have been present since some early epoch in their formation history. A further possibility to explain the metallicity gradients in cooling-flow clusters is that most of the metals in cluster cores may reside in large grains. The lifetime of grains of radius $a\\mu$m to sputtering in hot gas of density $n$ is $\\sim 2\\times 10^6 a/n\\yr$ (Draine \\& Salpeter 1979). Provided that individual grains exceed 10$\\mu$m in radius, they should survive for a Hubble time or longer throughout NCF objects and beyond the cooling radius in clusters with cooling flows. Within cooling flows, the density rises inward so the grains are increasingly sputtered, releasing the metals into the gas phase which thus becomes increasingly metal rich toward the cluster centre. Such a model requires that about half the metals are originally injected into the ICM as large grains. We note that large grains are inferred to have formed and to carry most of the iron in the expanding remnant of SN1987A (Colgan \\etal 1994). Within this model the typical mean metallicity in the {\\it core} of a cluster would be 0.5--1.0 solar. The grains would also provide a possible source for the dust inferred in the central optical nebulosities in cooling flows (Fabian, Johnstone \\& Daines 1994b; Voit \\& Donahue 1995; Allen \\etal 1995) and may be related to the excess soft X-ray absorption observed in cooling flows (\\eg Allen \\& Fabian 1997). \\subsection{The effects of cluster evolution} Within standard formation scenarios, clusters that form earlier are expected to have higher luminosities for a given temperature (since the gas density is higher at earlier formation epochs). Scharf \\& Mushotzky (1997) presented results from a single-temperature analysis of $\\sim 30$ clusters observed with ASCA and demonstrated a positive correlation between the amplitude, $A_{\\rm LT}$, of the $L_{\\rm X}-T_{\\rm X}$ relation (where they define $L_{\\rm X} = A_{\\rm LT}T_{\\rm X}^3$) and the mean (emission-weighted) metallicity. They suggest that the origin of this correlation is that clusters with larger values of $A_{\\rm LT}$ formed earlier and were better able to hold on to the metal-enriched gas expelled from their galaxies. Assuming that metals cannot be lost from the cluster potentials without a corresponding decrease in the X-ray gas mass, clusters with lower $A_{\\rm LT}$ should have lower baryon fractions by about a factor of two, given the observed range in metallicities. The Table lists the baryon fractions at a radius of 500 kpc in the clusters, determined with the appropriate spectral models (spectral model C for the CF systems and model A for the NCFs). We see that the mean baryon fractions for the subsamples of CF and NCF clusters differ by only $\\sim 30$ per cent, and thus that relatively little gas has escaped the potentials of the NCF clusters relative to the CF systems. (We note, however, that the distributions of baryon fraction values for the CF and NCF clusters are different. The application of a Student's t-test shows the mean values for the two subsamples to differ at $>99$ per cent confidence. The NCF clusters also have a significantly smaller dispersion in baryon fractions than the CF systems, and a Kolmogorov-Smirnov test shows the two subsamples to be drawn from different populations at $>99$ per cent significance.) If spectral model B rather than model A is used for the NCF systems ({\\it i.e.} if the absorbing column density is included as a free parameter in the spectral analysis of the NCF systems) the mean baryon fraction for these clusters rises to 0.16, improving the agreement with the CF clusters. It is important to note that the clusters studied in this Letter are amongst the most X-ray luminous and by implication most-massive clusters known. Material expelled from their member galaxies, even when part of an early low-mass subclump, is unlikely to escape from the total cluster potential. The accompanying Letter (Allen \\& Fabian 1998) discusses the impact of cooling flows on the $L_{\\rm Bol}-T_{\\rm X}$ relation for clusters. Allen (1998) discusses the effects of cooling flows on X-ray mass measurements. The new data presented in this Letter reveal a link between cooling flows and metallicity measurements and suggest the presence of metallicity gradients in clusters with cooling flows. Such effects must be accounted for before attempting to determine cosmological parameters from X-ray observations of clusters." + }, + "9802/astro-ph9802263_arXiv.txt": { + "abstract": "We describe a system in use at the Lick Observatory 1-m Nickel telescope for near-simultaneous imaging at optical and near-infrared wavelengths. The combined availability of a CCD and a NICMOS-3 camera makes the system well-suited for photometric monitoring from 0.5-2.2$\\micron$ of a variety of astrophysical objects. Our science program thus far has concentrated on studying variability trends in young stellar objects. ", + "introduction": "There is broad need in astronomy for simultaneous optical-infrared photometric capabilities which enable temporal monitoring. Examples include recording supernova light curves, monitoring variability of active galactic nuclei, and measuring the spectral energy distribution of young stars. Because supernovae are standard candles, their light curves are important and have maximum utility if well sampled and cover multiple wavelengths. Dense temporal sampling is valuable because this helps to establish the epoch of maximum light, a task that is simplified for type Ia events if $J$-band data are available, and some immunity to the effects of dust can be achieved if observations extend into the $K$-band where extinction is an order of magnitude lower than at visible wavelengths (Elias \\& Frogel 1981). Variability of AGN such as quasars and Seyfert nuclei provides an opportunity to investigate the nature of the central engine, and the distribution and properties of dust grains via reverberation mapping (e.g., Nelson 1996). Part of the data for Nelson's (1996) study came from a dual optical-IR imager on the UCLA 24-inch Telescope (Nelson et al. 1997). Another example of a simultaneous optical-IR system is the $IJK$ camera used for the DENIS survey (Copet et al. 1997). Our primary interest thus far has been to pursue an optical-infrared photometric monitoring program to investigate and characterize variability trends in the spectral energy distributions (SEDs) of young stars. Circumstellar accretion disk models matched to the observed SEDs of young pre-main sequence stars are used to attempt understanding of the complex accretion processes in these systems. A major uncertainty in the SED modeling procedure, however, is the continuum variability which occurs at all wavelengths out to at least 3.5 $\\mu$m. Changes in flux of several tens to several hundreds of percent occur on time-scales of a few hours to a few days, depending on the source and on the wavelength regime. There are several possible causes for the observed monochromatic and color variability: 1) variable stellar flux; 2) variable extinction; and 3) variable properties of the circumstellar material. Sufficient optical and near-infrared monitoring has been carried out for a few sources so that we are able to investigate their complicated behavior in color-color and color-magnitude diagrams ({\\it c.f.} Skrutskie et al. 1996; Herbst, Herbst, \\& Grossman 1994). These patterns seem inconsistent with simple variability explanations like changes in intrinsic stellar flux or line-of-sight obscuration (although these causes can not yet be discounted entirely). Thus it may be that the observed variability trends are due to changes in the physical properties of the hot circumstellar gas and dust, e.g., temperature, opacity, or geometry variations taking place in accretion columns, in the inner disk, or in the stellar/disk wind. Testing this hypothesis requires a more extensive photometric dataset than currently exists. Relevant monitoring time-scales might encompass: the dynamical time for a wind (R$_*$ / v$_{wind} \\sim$ hours), the free-fall time from the inner disk to the stellar surface ($\\sqrt{r^3/GM}\\sim$ tens of hours), the dynamical time for the disk ($1/\\Omega = \\sqrt{r^3/GM}\\sim$ tens of hours to years for r = 0.1-1AU), the dust destruction / formation time ($\\sim$ years), and the viscous accretion time-scale ($\\alpha^{-1}~r^2/h^2~\\tau_{rot}\\gtrsim$ tens of years). Both scientific and practical considerations shaped our plans for implementing simultaneous optical-infrared measurements at the University of California's Lick Observatory on Mt Hamilton. Our need to image complex fields with multiple and sometimes extended sources ruled out a single-channel photometer. However, as cost precluded development of new cameras, combining existing instruments to meet our broadband imaging requirements presented an attractive design strategy. This approach also permitted us to take advantage of existing electronic and software controls, and to adopt existing mechanical telescope interfaces with only minor modifications. Since the science goals require intensive photometric monitoring of relatively bright targets, we concentrated our instrument development effort on the 1-m Nickel telescope, rather than the heavily over-subscribed 3-m Shane telescope (though, in principle, VIRIS could be used with either). We thus arrived at the option of combining Lick's 1-m facility CCD camera and the facility IR camera, used at both the 1- and 3-m telescopes. ", + "conclusions": "We have designed and constructed an efficient system for near-simultaneous optical and IR imaging and used it to obtain accurate $VRIJHK$ photometry in a program to monitor the variability of T-Tauri stars. VIRIS demonstrates the feasibility of combining existing optical and IR instruments using an opto-mechanical interface to permit simultaneous and near-simultaneous operation. We have shown that the details of the interface are very simple, and likely could be replicated on other small telescope to permit them to tackle similar scientific programs that require panchromatic observations." + }, + "9802/astro-ph9802113_arXiv.txt": { + "abstract": "Stellar nucleosynthesis is the corner-stone of many astrophysical problems. Its understanding, which can be tested by countless observations, leads to insights into the stellar structure and evolution, and provides crucial clues to the physics of the galaxies and of the universe. Precise answers can be given to the questions ``When, where and how the chemical elements are synthesized in stars?\", and will be summarized in this paper for what concerns stars of low and intermediate mass. However, in spite of the observational confirmation of many predictions, important and complementary data reveal some weaknesses in the theory of stellar physics. In particular, the mixing processes of chemicals inside the stars, as well as the mass loss and ejection mechanisms, are poorly known. Their rudimentary treatment imposes the use of parameters which strongly influence the predicted chemical yields. In our discussion, we will underline the future developments that may lead to quantitative changes in the predictions of chemicals production by low and intermediate mass stars. ", + "introduction": "We restrict our discussion to the case of single low and intermediate mass stars (LIMS), which are those that develop an electron-degenerate carbon-oxygen core after the center He-exhaustion and do not proceed through the carbon and heavier elements burning phases (The case of massive stars is treated in this volume by N.~Prantzos). This definition places an upper limit of 6--8\\,M$_{\\odot}$ (see Maeder \\& Meynet 1989 for a discussion on this value which depends mainly on the mixing prescription and mass loss rates used in the evolutionary models). \\subsection{The Main Trends} The nucleosynthesis occuring during the evolution of LIM stars is well known. For what concerns the stellar surface abundances and the chemical enrichment of the interstellar matter, the only interesting nucleosynthesis process in these stars up to the AGB phase is hydrogen burning (the products of central helium burning remain indeed trapped into the white dwarf). Nuclear reactions involving elements from D to C begin to occur on the pre-main sequence for temperatures higher than 10$^6$\\,K. During this phase, D burns to $^3$He; $^6$Li and $^7$Li are destroyed by p-capture, and their abundances decrease at the surface of stars with masses lower than $\\sim$ 1\\,M$_{\\odot}$. On the main sequence (MS), hydrogen burning occurs in a radiative core and is dominated by the pp-reactions in stars with masses lower than $\\sim$ 1.1\\,M$_{\\odot}$, while the CNO bi-cycle dominates in the convective core of more massive stars. Partial pp-chain burning builds up a peak of $^3$He. Slightly further in, $^{12}$C and $^{16}$O are partially converted into $^{14}$N. A peak in the $^{13}$C abundance develops. Below, $^{18}$O is burned to $^{15}$N, and $^{17}$O is slightly enhanced from partial burning of $^{16}$O (a peak of $^{17}$O is present). The chemical profiles at the end of the MS depend both on mass and metallicity. After the central H exhaustion, nuclear energy is released by the hydrogen burning shell (HBS) that surrounds the He-core which becomes electron degenerate in low mass stars (LMS, i.e., with masses lower than 2-2.3M\\,$_{\\odot}$). The star expands and the convective envelope penetrates deeply inwards until $\\sim 70 \\%$ of the total mass is convective, reaching the region where central H-burning occured. The first dredge-up (1Dup) leads to the first important modification of the surface abundances. While the star is on the horizontal branch, the convective core where helium burns is surrounded by a thin HBS. Following core helium exhaustion, the structure readjusts to shell burning (HeBS, HBS), and the star strongly expands while it enters the early-AGB phase. The envelope penetrates in the deep layers. In stars with masses higher than $\\sim$ 4M\\,$_{\\odot}$, a second dredge-up (2Dup) mixes to the surface the products of complete H-burning (contrary to the 1Dup, the HBS extinguishes). Soon after, the first thermal pulse (TP) occurs (see \\S 2.1). \\subsection{The First and Second Dredge-Up. Theoretical Predictions} During the 1 and 2Dup, convective mixing and induced dilution modify the surface abundances of the elements altered by H-burning. The extent of these changes depends both on stellar mass and metallicity (Z): For stars with masses higher than 2-2.3M\\,$_{\\odot}$, the 1Dup is almost negligible at low metallicity, and its importance increases with Z. The effect of the 2Dup in lower mass stars is negligible whatever the metallicity. The predictions obtained recently by different groups agree quantitatively well (Schaller et al. 1992; Dearborn 1992; Bressan et al. 1993; El Eid 1994; Forestini \\& Charbonnel 1997 (hereafter FC97); Boothroyd \\& Sackman 1997 for the most recent papers), except for a few features. Basically, one can summarize as follows: In LIM stars, the surface mass fraction of $^3$He increases by factors 2 to 6. Once in the convective layers of the red giant, $^3$He is preserved because of the too cool temperature in these regions. Since $^7$Li is burned on the main sequence in the regions where the temperature is higher than about $2.5 \\cdot 10^6$\\,K, the surface abundance of this element strongly decreases after the 1Dup. The total carbon abundance decreases by approximately 30$\\%$; there is no carbon depletion for $<$1M\\,$_{\\odot}$ stars, and less than a factor of 2 at higher masses. From an initial value of the order of 90, the carbon isotopic ratio decreases down to about 20-30 due to the 1Dup; since after central HeB, almost all the CNO isotopes have been converted into $^{14}$N, the 2Dup has little effect on the $^{12}$C/$^{13}$C ratio. The $^{14}$N abundance increases by about 80$\\%$ and the $^{12}$C/$^{14}$N decreases. The $^{16}$O/$^{17}$O ratio does not change for stars with masses lower than 1M\\,$_{\\odot}$, but it decreases when the convective envelope reaches partially the $^{17}$O pocket (in stars with masses between 1 and 2M\\,$_{\\odot}$) or completely engulfes it (in more massive stars; there, uncertainties in the $^{17}$O-destruction rates affect the predictions for the final $^{16}$O/$^{17}$O ratio). The very slight rise of the $^{16}$O/$^{18}$O ratio is an increasing function of the stellar mass. The total abundance of oxygen and of all heavier elements is not affected by the 1 and 2Dup. \\subsection{Comparisons with the Observations} Many observations exist to which one can confront the predictions described above. For what concerns stars with masses higher than about 2M\\,$_{\\odot}$, no major conflict appears. The theoretical post dredge-up values of the carbon isotopic ratio are slightly lower but in agreement with the observations in galactic cluster giants (Gilroy 1989). Red giants present \\chem{O}{16}/\\chem{O}{17} ratios between 300 to 1000 and \\chem{O}{16}/\\chem{O}{18} ratios in the range 400 to 600 (Harris et al. 1988; Smith \\& Lambert 1990a), as predicted. The case of lower mass stars is however problematic (see Charbonnel et al. 1998, hereafter CBW98, for references). In most of the LM evolved stars, the observed conversion of $^{12}$C to $^{13}$C and $^{14}$N greatly exceeds the levels expected from standard stellar models. The $\\rm ^{12}C/^{13}C$ ratio even reaches the near-equilibrium value in many Population II RGB stars. This problem also occurs, but to a somewhat lower extent, in evolved stars belonging to open clusters with turnoff masses lower than 2\\,M$_{\\odot}$. In halo giants, the lithium abundance continues to decrease after the completion of the first dredge-up. A continuous decline in carbon abundance with increasing stellar luminosity along the RGB is observed in globular clusters such as M92, M3 and M13, M15, NGC 6397, NGC 6752 and M4. In some globular clusters (M92, M15, M13, $\\omega$ Cen), giants exhibit evidence in their atmospheres for O$\\rightarrow$N processed material. In addition to the O versus N anticorrelation, the existence of Na and Al versus N correlations and Na and Al versus O anticorrelations in a large number of globular cluster red giants has been clearly confirmed. \\subsection{Towards More Complete Stellar Models. Consequences of an Extra-Mixing on the RGB} None of the behaviours described above is predicted by the standard stellar theory. These observations suggest that, while they evolve on the RGB, LM stars undergo an extra-mixing in the region situated between the hydrogen burning shell (where the material is processed through the CN-cycle and possibly the ON-cycle) and the deep convective envelope. This extra-mixing, which appears to be efficient only after the luminosity-function bump (Gilroy \\& Brown 1991; Charbonnel 1994; CBW98), adds to the standard first dredge-up to modify the surface abundances. Recently, different groups have simulated an extra-mixing process in order to reproduce the CNO abundances in RGB stars (Denissenkov \\& Weiss 1995; Wasserburg et al. 1995; Boothroyd \\& Sackman 1997). Some authors attempted to relate the extra-mixing with physical processes, among which rotation and mass loss seems to be the most promising (Sweigart \\& Mengel 1979; Charbonnel 1995). In any case, these ``non-standard\" models (related or not to a physical process), make a common crucial prediction : The mechanism(s) which is (are) responsible for the chemical anomalies on the RGB must lead to the destruction of $^3$He by a large factor in the bulk of the envelope material (Hogan 1995; Charbonnel 1995). This result strongly modifies the actual contribution of LM stars to the galactic evolution of this element (see Tosi 1996; Charbonnel 1997). It prevents its overproduction and helps understanding the recent measurements of $^3$He in the local interstellar cloud (Gloeckler \\& Geiss 1996), in galactic H\\,{\\sc ii} regions (Balser et al. 1994) and in planetary nebulae (Rood et al. 1992; Balser et al. 1997). The physics of the extra-mixing process in LMS has to be better understood. Detailed simulations, with a consistent treatment of the transport of matter and angular momentum, have to be carried out for different stellar masses and metallicities, and various mass loss and rotation histories. The impact of this process on the behavior of various elements in RGB stars (C $\\searrow$, Na $\\nearrow$, O $\\searrow$, Al $\\nearrow$, Mg O $\\searrow$) and on the precise yields of $^3$He has to be investigated in details. Consequences for the energy production in the HBS and for the HB morphology may not be neglected (Sweigart 1997). ", + "conclusions": "" + }, + "9802/astro-ph9802325_arXiv.txt": { + "abstract": "Spiral density wave theories demand that grand design spiral structure be bounded, at most, between the inner and outer Lindblad resonances of the spiral pattern. The corotation resonance lies between the outer and inner Lindblad resonances. The locations of the resonances are at radii whose ratios to each other are rather independent of the shape of the rotation curve. The measured ratio of outer to inner extent of spiral structure for a given spiral galaxy can be compared to the standard ratio of corotation to inner Lindblad resonance radius. In the case that the measured ratio far exceeds the standard ratio, it is likely that the corotation resonance is within the bright optical disk. Studying such galaxies can teach us how the action of resonances sculpts the appearance of spiral disks. This paper reports observations of 140 disk galaxies, leading to resonance ratio tests for 109 qualified spirals. It lists candidates that have a good chance of having the corotation resonance radius within the bright optical disk. ", + "introduction": "Disk galaxies exhibit many designs: bars, rings, and spiral structure. The structure and placement of rings and bars are widely believed, on numerical and analytical grounds, to be governed by resonance. A resonance is a rational relationship between two frequencies. Inner and outer rings form at radii that are resonant with the angular frequency of the bar pattern (Schwarz \\markcite{Schw81}1981; Buta \\markcite{Buta86}1986; Buta \\& Crocker \\markcite{Buta91}1991). A strong bar ends near its corotation resonance (Contopoulos \\markcite{Cont80}1980), referred to hereafter as ``CR''. Many disk galaxies show spiral patterns, some of which are beautiful, two-armed spirals. What effect do resonances have on the appearance of spiral patterns? First consider barred spirals. It was once widely believed that the bar was capable of ``driving'' spiral structure (e.g., van Albada \\& Roberts \\markcite{vanA81}1981) or of providing a source of disturbances that might be swing amplified into spiral structure (Toomre \\markcite{Toom77}1977). Then the bar and the spiral patterns would be resonantly related. But Sellwood \\& Sparke (\\markcite{Sell88}1988) have suggested that spiral patterns need not be resonantly related to bars. Kinematic data for few galaxies have been analyzed to see if the resonant link exists between bars and spirals. Bar and spiral are part of the same pattern in both NGC1365 and NGC4321, for example, but the transition from bar to spiral is at the CR in NGC1365 (Lindblad, Lindblad, \\& Athanassoula \\markcite{Lind96}1996) while it is at the inner 4:1 resonance in NGC4321 (Elmegreen, Elmegreen, \\& Montenegro \\markcite{Elme92}1992; Sempere et al. \\markcite{Semp95}1995; Canzian \\& Allen \\markcite{Canz97}1997). Apart from these two galaxies, there are differences between strongly and weakly barred spirals to consider, as well as differences between flat and exponential bars. Spiral patterns exist in disks without bars, too. Spiral density waves are constrained to exist between the inner and outer Lindblad resonances, hereafter ``ILR'' and ``OLR''. Aside from this constraint, the role of resonances in shaping pure spiral morphology is uncertain. Many signatures of resonance structures have been suggested and used with varying success. Dust lanes may jump to the other sides of spiral arms across the CR (Roberts \\markcite{Robe69}1969). Star formation may be less efficient within spiral arms near CR (seen in NGC0628 and NGC3992 by Cepa \\& Beckman \\markcite{Cepa90}1990 and in M51 by Knapen et al. \\markcite{Knap92}1992 but not in NGC4321 by Knapen et al. \\markcite{Knap96}1996). Enhanced star formation may occur in the interarm region at the CR, and at the inner 4:1 resonance there may be brightness minima in the arms and spurs should erupt in the interarm region (Elmegreen et al. \\markcite{Elme89}1989, \\markcite{Elm92}1992). Not all of the above prescriptions work in all cases. Tests of the above prescriptions can follow only when clear determinations of resonance locations have been made for many spirals. Confirmation of resonance locations requires kinematic data. Such data derive from emission line mapping, which is difficult and time consuming. It would be more rewarding to study galaxies for which interesting resonance locations were known beforehand to be at easily studied radii in the bright disk. The CR is one of the most interesting resonances to study. This paper presents a test to identify galaxies whose CR is favorably placed for study (within the bright disk). The test requires only moderately deep imaging in typical seeing. No rotation curve data are necessary. The test is not completely determinate: galaxies with CR within the bright disk can be missed. This is the price paid for operating in ignorance of the rotation curve shape and spiral pattern speed. However, the test is reliable: when a suitable galaxy is identified, there is a good chance that its CR is, indeed, within the bright disk. The basis of the test to identify galaxies with well-placed CR is presented in \\S\\ref{sec_extensive}. The section also investigates how much effect the shape of the rotation curve and the spiral pattern speed have on the outcome of the test. \\S\\ref{sec_observations} outlines the observations. The analysis of the imaging is described in \\S\\ref{sec_analysis}. This includes not only a description of the measurements made but also some surface photometry to estimate how deep the imaging reaches. The results are presented in \\S\\ref{sec_discussion}, where promising galaxies are described and where the different characteristics of barred and pure spirals are discussed. ", + "conclusions": "This paper describes an observational project that examined 140 spiral galaxies, measuring the visual extent of their spiral structure. The purpose of the project was to find galaxies for which the corotation resonance of the two-armed grand design spiral structure was likely to be within the bright optical disk. The likelihood that the corotation resonance is within the optical disk can be assessed by comparing the measured ratio of outer to inner spiral extent with the value of the canonical ratio (equation~\\ref{eq_ilrcr}). The test using the canonical ratio is not conclusive because of variations due to rotation curve shape. It is also possible that there is no inner Lindblad resonance. To counter our ignorance of the rotation curve, a test with a more conservative value of the ratio than equation~\\ref{eq_ilrcr} should be used (here, an extra factor of two). A few galaxies were found that have large spiral extent and whose corotation resonance radius is very likely to be within the bright optical disk. Among the best candidates for kinematic study are NGC5829, NGC6907, IC0211, NGC1042, and NGC0132. It was incidentally found that barred and ringed spirals statistically have more limited spiral extent than pure spirals. This observation can be explained if many bars end at their inner 4:1 or corotation resonance and if the bar and spiral are part of the same pattern. The location of inner rings at the inner 4:1 resonance also limits spiral extent, but the location of outer rings at the outer Lindblad resonance places no extra constraint on spiral extent." + }, + "9802/astro-ph9802107_arXiv.txt": { + "abstract": "s{Number counts and redshift distribution of gravitational arcs are computed in the field of massive clusters of galaxies to probe the universe at high redshift. Using an accurate modelling for the cluster mass distribution and a model for the spectrophotometric evolution of galaxies, the redshift distribution of gravitational arclets is computed in the field of cluster Abell 2218 and in the Hubble Deep Field where a cluster is artificially located. Counts are very well reproduced in the $B$ band but an important population appears at high redshift which is not seen in deep spectroscopic surveys. Unfortunately, the very high sensitivity of the counts with respect to the model for galaxy evolution and to the mass distribution prevents from estimating the cosmological parameters with arcs statistics. Future works have to concentrate on high redshift clusters and take advantage of objects with smaller distortions.} ", + "introduction": " ", + "conclusions": "It is now obvious that arcs statistics can only be performed with mass distributions including substructures, too simple potentials should be avoided. Number counts of arclets are in good agreement with observations in the $B$ band and the redshift distribution of arclets at $z \\le 1.0 $ is correctly predicted by the model. The important population of arclets expected at $z \\ge 1.0 $, which is not observed in spectroscopic surveys, is highly dependent on the modelling of high redshift ellipticals and the role of dust absorption in the rest frame UV. Unfortunately, the geometry of the universe cannot be determined with arcs statistics because of too important uncertainties in clusters mass distributions and in the model for galaxy evolution. Geometrical effects have now to be investigated with clusters at redshift greater than 0.5 considering also the weak lensing regime. High-$z$ lens efficiency relates more directly to cosmology and counts of all the background objects in a cluster field enable accurate measurement of the magnification bias that probes the high redshift population. \\vskip -0.1cm" + }, + "9802/astro-ph9802277_arXiv.txt": { + "abstract": "I examine the role of dust grains in determining the structure of steady, cold, oblique C-type shocks in dense molecular gas. Gas pressure, the inertia of the charged components, and changes in ionization are neglected. The grain charge and rate coefficients for electron-neutral and grain-neutral elastic scattering are assumed constant at values appropriate to the shock interior. An MRN size distribution is accounted for by estimating an effective grain abundance and Hall parameter for single-size grains. A one-parameter family of intermediate shocks exists for each shock speed \\( v_s \\) between the intermediate signal speed \\(v_A {\\rm cos} \\theta \\) and \\(\\sqrt{2} v_A {\\rm cot} \\theta\\), where \\( v_A \\) is the pre-shock Alfv\\'en speed and \\(\\theta\\) is the angle between the pre-shock magnetic field and the normal to the shock front. In addition, there is a unique fast shock for each \\( v_s > v_A \\). If the pre-shock density \\( \\nh \\ga 10^5 \\percc \\) and the pre-shock magnetic field satisfies \\( B(\\mathrm{mG}) /\\nh(10^5\\percc ) \\la 1 \\) grains are partially decoupled from the magnetic field and the field and velocity components within fast shocks do not lie in the plane containing the pre-shock field and the shock normal. The resulting shock structure is significantly thinner than in models that do not take this into account. Existing models systematically underestimate the grain-neutral drift speed and the heating rate within the shock front. At densities in excess of \\( 10^8 \\percc \\) these effects may be reduced by the nearly-equal abundances of positive and negative grains. ", + "introduction": "The structure of shock waves in molecular clouds is determined by the coupling between the magnetic field and the weakly ionized pre-shock gas. For shock speeds below 40--50 \\( \\kms \\) (McKee, Chernoff \\& Hollenbach 1984) Lorentz forces within the shock front push the charged species through the neutrals, and the resulting collisions accelerate, compress, and heat the ambient gas. This process is slow because the charged species are rare, so the gas is able to radiate away a significant fraction of the heat while still within the shock front (Mullan 1971; Draine 1980), primarily through the emission of radiation in molecular rotational and vibrational transitions. Shocks of this nature are denoted `C-type' (Draine 1980), or `\\cstar-type' if the gas is decelerated through a sonic point in the reference frame comoving with the shock front (Chernoff 1987, Roberge \\& Draine 1990). At higher shock speeds, the molecular coolants are dissociated, cooling can no longer keep pace with heating, and the gas pressure becomes dynamically significant. A thin, viscous sub-shock forms within the front, and the shock is termed `J-type' (Draine 1980). C-type shocks efficiently convert the heat produced by collisions within the shock front into molecular line emission, and are responsible for much of the intense infrared H\\( _2 \\) and CO line emission observed towards the Orion-KL region (Draine \\& Roberge 1982; Chernoff, McKee \\& Hollenbach 1982; Smith \\& Brand 1990; Smith, Brand \\& Moorhouse 1991; Chrysostomou et al. 1997). The warm molecular environment within C-shocks can drive molecular chemistry such as the conversion of atomic oxygen into OH and then into water (Draine, Roberge \\& Dalgarno 1983; Kaufman \\& Neufeld 1996a,b). In addition, the large drift speeds can drive exothermic ion-neutral chemistry within the shock front (Flower, Pineau des F\\^orets \\& Hartquist 1985; Pineau des F\\^orets, Flower \\& Hartquist 1986; Draine \\& Katz 1986a,b). Shock models have generally been forced to be `coplanar', that is, the drift velocities and magnetic field within the shock front are assumed to lie in the `shock plane', the plane containing the pre-shock magnetic field and the shock normal (see Fig. 1). This holds for fast\\footnote{`Fast' and `slow' have occasionally been used in the astrophysical literature to denote J-type and C-type shocks respectively. In this paper I use `fast' and `slow' in the traditional MHD sense, referring to the fast and slow signal speeds.} shocks if the charged species are tied to the field lines, but charged grains drift obliquely to the shock plane if they become partially decoupled from the field lines by collisions with the neutrals (Draine 1980). However, for the sake of simplicity, the component of the grain drag force on the neutrals perpendicular to the shock plane has usually been supressed (e.g. Draine et al. 1983; Wardle \\& Draine 1987). This approximation breaks down for high pre-shock densities, partly because the fractional ionization of the gas decreases, permitting the grains to make a larger fractional contribution to the drag on the neutral gas, but also because the grains become poorly coupled to the magnetic field and the grain drag vector is further tilted out of the shock plane. This motivated calculations of the structure of grain-dominated C-shocks in which the vector components perpendicular to the shock plane were retained. Pilipp, Hartquist \\& Havnes (1990) showed that the thickness of perpendicular shocks was significantly changed for pre-shock densities in excess of \\(10^7 \\percc\\). In this case, the symmetries of the geometry constrain the magnetic field direction to be the same throughout the shock front, even though the drift velocities need not lie within the shock plane. Oblique shocks introduce an important extra degree of freedom, in that \\( \\Bp \\), the magnetic field component perpendicular to the shock normal, may rotate within the shock front before returning to the shock plane downstream. Oblique shocks were studied by Pilipp \\& Hartquist (1994), who found solutions for shock speeds below a few \\(\\kms\\) in which \\( \\Bp \\) rotates by 180\\degr (in either sense) within the shock front. Pilipp \\& Hartquist were unable to find solutions for shock velocities in excess of a few \\(\\kms\\), and suggested that steady shock solutions may not exist in this regime. This is disconcerting, as shocks with speeds in excess of 20 \\(\\kms\\) are required to model the observed emission from a variety of sources. In addition, the development of a sound theoretical understanding of magnetohydrodynamics of dense, dusty gas is important for scenarios for the formation and collapse of molecular cloud cores and the formation of protostellar discs (e.g. Nishi, Nakano \\& Umebayashi 1991; Ciolek \\& Mouschovias 1993; Neufeld \\& Hollenbach 1994; Li \\& McKee 1996; see also the review by Hartquist, Pilipp \\& Havnes 1997). The behaviour of \\( \\Bp \\) in the solutions obtained by Pilipp \\& Hartquist is fundamentally different to the earlier coplanar oblique shock models (Wardle \\& Draine 1987; Wardle 1991a; Smith 1992), in which the pre-shock and post-shock \\( \\Bp \\) vectors are parallel rather than antiparallel. The upstream and downstream states across magnetohydrodynamic (MHD) shocks are related by jump conditions that are independent of the detailed nature and behaviour of the fluid, and therefore of the actual shock structure. In particular, the transverse magnetic field increases, changes sign, or decreases across the fast, intermediate, or slow shocks respectively (e.g. Cowling 1976; Kennel, Blandford \\& Coppi 1989). This implies that the solutions found by Pilipp \\& Hartquist are C-type \\emph{intermediate} shocks, whereas previously studied C-shock solutions are \\emph{fast} shocks. The jump conditions show that the downstream state for intermediate shocks is unphysical for shock speeds a little above the Alfv\\'en speed, and this explains the breakdown of the solutions for mildly super-Alfv\\'enic shocks. Further, for a given shock speed, Pilipp \\& Hartquist found a one-parameter family of shock solutions rather than the single shock solution found in previous studies, and noted that this has also been found to be the case for intermediate shocks in resistive MHD (Wu 1988a; Kennel, Blandford \\& Wu 1990). Having identified these solutions as intermediate shocks, there still remains the important issue of the existence of fast shock solutions. On grounds of physical continuity, one expects that adding grains to a coplanar, fast C-type shock should distort the structure out of the shock plane rather than do away with it entirely. Indeed, there are hints that steady shock solutions may exist. Following standard practice in C-shock modelling, Pilipp \\& Hartquist derived a set of ODE's describing the shock structure, and integrated them from a perturbed upstream state towards the shock front. There is, however, a difference in the number of free parameters describing the initial perturbed state. In the coplanar case there is one free parameter -- the amplitude of the initial perturbation in the pre-shock fluid (Draine 1980; Wardle \\& Draine 1987). Physically, this specifies the point in the shock precursor at which the integration towards the shock front begins, so there is a single shock solution for a given shock speed, the fast shock. Dropping the assumption of coplanarity introduces a second parameter, which can be regarded as the degree of rotation of the perturbation out of the shock plane at the initial point. The freedom introduced by the new parameter potentially allows a family of solutions to exist (the intermediate shocks), and Pilipp \\& Hartquist found that there is a critical value at which the sense of rotation of \\( \\Bp \\) across the intermediate shock solutions changes sign. This suggests that the fast shock solution corresponds to this critical value, but that the solution cannot be found by integration from the perturbed upstream state because of finite numerical precision. This paper demonstrates that fast non-coplanar C-shock do exist. As the issues addressed here are fundamental in nature, I keep the formulation of the multifluid shock problem, described in the next section, as simple as possible. In particular, I assume that the molecular gas consists of four distinct species: neutrals, positive ions, electrons or PAHs, and negatively charged grains. Gas pressure is neglected, and ionization balance and chemistry are also ignored. These simplifications allow all of the physical quantities in the shock to be expressed in terms of the components, \\(B_x\\) and \\(B_y\\), of \\( \\Bp \\). Thus a shock solution can be represented by a plot of \\(B_y(z)\\) {\\it vs.} \\(B_x(z)\\), where \\( z \\) is the coordinate along the shock normal. This two-dimensional phase space, explored in \\S\\ref{sec-analysis}, is a powerful tool for examining the set of shock solutions given the pre-shock conditions and shock speed. In \\S\\ref{sec-results} I show that both fast and intermediate shock solutions exist for low shock speeds, and that the topology of the phase space prevents the fast shock solution from being found by integration from upstream to downstream. At higher speeds the intermediate solutions become unphysical, but the fast solutions persist. I show that marginally coupled grains dominate the grain drag, and that their effect on the shock structure is dramatic, and examine the effects of an MRN grain-size distribution. The implications of these results are discussed in \\S\\ref{sec:discussion}, and a summary is presented in \\S\\ref{sec:summary}. ", + "conclusions": "\\label{sec:discussion} The results presented in \\S\\ref{sec-results} confirm that both intermediate and fast C-type shock solutions exist. A one-parameter family of intermediate shocks exists for each shock speed between the intermediate speed \\( v_A \\cos\\theta \\) and \\( \\sqrt{2} v_A \\cot \\theta \\); at higher speeds the downstream states become unphysical. A unique fast shock exists for each shock speed in excess of the fast speed \\( v_{A} \\). The relationship between these two classes of solution is summarised by the trajectories in the \\( B_{x} \\)--\\( B_{y} \\) phase-space diagram. The fast shock trajectory runs along the \\( B_{x} \\) axis if the charged species are tied to the magnetic field lines, or if positive and negatively charged species have identical abundances and Hall parameters. In this case, integration of the fast shock solution away from the upstream state is straightforward. The direction of the trajectory leaving the upstream stationary point (i.e. along the \\(B_x\\) axis) is known, and the velocity and magnetic field vectors in the fast shock solution are confined to the \\(x\\)-\\(z\\) plane, so the \\( y \\) components can be discarded. Integration can be started by stepping off U towards F, and then integrating the trajectory (e.g. Draine 1980; Wardle \\& Draine 1987). The symmetry of phase space under reflection in the \\( B_x \\) axis is destroyed by grains, which being partially decoupled from the magnetic field and preferentially carrying a negative charge, introduce a handedness into the fluid. The fast shock trajectory no longer runs along the \\( B_{x} \\) axis, and this distortion provides an opportunity for numerical mixing of the fast solution with the neighbouring divergent intermediate shock trajectories during integration. Any attempted integration from U to F is doomed. To begin with, there is no {\\it a-priori} means of determining the direction for the initial step off the point U. In principle this problem could be circumvented by using a shooting method to adjust the starting direction until the integration successfully reached F. However, because of finite numerical precision, any attempted integration into F along the fast trajectory will step onto a neighbouring intermediate shock trajectory and eventually be forced away from the fast trajectory and towards I (c.f. integration towards sonic points). This behaviour prevented Pilipp \\& Hartquist (1994), who started their integrations at U and treated the direction of the first step away from U as a free parameter, from finding the fast shock solutions. At low shock speeds, they discovered the family of intermediate shocks but could not find the fast shock, noting instead that the sense of rotation of \\( \\Bp \\) within the intermediate shocks changed at a critical direction of the initial step away from U. At higher speeds, for which the intermediate shock solutions become unphysical, Pilipp \\& Hartquist were unable to find any acceptable solutions. A realistic treatment of the ionization balance is required for a quantitative assessment of the role of grains in C-type shocks. The ionization balance is particularly important when \\( v_z \\) for any charged species approaches zero and the number density of that species diverges in the constant-flux approximation (eq. \\ref{eq-jcontinuity}) used here. This does not affect the divergence of trajectories from the fast trajectory unless the fast trajectory itself approaches this limit (e.g. the trajectory shown in Fig. \\ref{fig-fast-phase-space} for the shock displayed in Fig. \\ref{fig-t13-struc}). Some conclusions are robust to changes in the microphysics. If the conductivity of the gas can be regarded as some unspecified function of local physical variables at any point within the shock front, so that shock solutions can still be represented as trajectories in the \\( B_y \\)--\\( B_x \\) phase-space, the topology of the phase space is determined by the nature of the fixed points. These were shown in \\S \\ref{subsec-classification} to be a source, sink and saddle independent of the conductivity tensor. Thus the integral curves always have the topology plotted in Figure \\ref{fig-int-phase-space}, apart from smooth distortions (as apparent in Fig \\ref{fig-t6-phase-space}) because of changes to the conductivity tensor. The intermediate-shock curves always run from U or D to I, and thus \\( v_{iz}=v_{ez}=0 \\) at some point if the shock speed is high enough for intermediate shocks to be excluded. The fast shock trajectory could in principle cross the \\( v_{iz}=v_{ez}=0 \\) locus if the ionization balance is able to distort the trajectory sufficiently, although it is apparently unable to do so in the fixed-flux approximation used here (see Fig. \\ref{fig-noPAH-phase-space}). In any case, when the fixed-flux condition is dropped, crossing this locus does not imply infinite charged-particle abundances as the increase is limited by recombinations. Intermediate shocks, however, will still be excluded as their downstream conditions are unphysical. A more realistic treatment of the physical processes occurring within the shock front is essential as the grain properties within the shock and the shock structure itself are intimately related. In particular, the neutral and electron temperatures are required to accurately calculate the Hall parameters for a given grain size. In addition, the grain-size distribution varies within the shock front as different sized grains will generally have different drift speeds, the smallest grains being compressed with the magnetic field, the largest grains with the neutrals. Incorporation of these effects increases the dimensionality of the phase space by introducing differential equations for the temperatures and relative abundances of different species. In principle this could alter the topology of the phase space, although this is unlikely to be the case in practice. There are, however, practical problems in finding the fast shock structure by integrating backwards from the downstream state if chemistry is coupled to the hydrodynamics (through e.g. changing the abundances of coolants or the degree of ionization in the shock), since one does not then know the post-shock conditions {\\it a-priori} (c.f. the discussion in Roberge \\& Draine (1990) regarding C\\(^*\\)- and J-shocks), and a time-consuming shooting method must be employed. Although the grain drag significantly modifies the shock structure, it is unclear whether this will produce gross differences in the total line emission from the shock front. The essential feature of C-shock structure is that it maintains a layer of hot molecular gas. For a given pre-shock density and shock speed, the same power is dissipated per unit area independently of the detailed structure of the shock front. The gas generally will be heated up to temperatures between 1000 and 2000 K at which point the local heating and cooling balance. Molecules will remain intact and the molecular line emission will not be greatly affected. The insensitivity of line ratios to details of shock structure is illustrated by recent numerical work (Stone 1997; Neufeld \\& Stone 1997; MacLow \\& Smith 1997), which follows the development and saturation of the instability in C-type shocks (Wardle 1990). Even though the neutral gas is collected into dense fingers within the shock front, the line ratios are generally unchanged by more than a factor of two (Neufeld \\& Stone 1997), although diagnostics can be found (MacLow \\& Smith 1997). Nevertheless, one might expect that the grains will affect the stability of C-shocks and also the speed at which shocks become J-type. Finally, it is worth pointing out that the theoretical study of C-shocks in molecular clouds will contribute to the general theory of MHD shock waves, and of intermediate shocks in particular (see the review by Wu 1995). Until recently intermediate shocks in MHD were thought to be unphysical because they appeared to lack sufficient freedom to adjust to slight perturbations in the upstream or downstream flow (see, e.g. Kantrowitz \\& Petschek 1965). The issue was reopened when intermediate shocks emerged as stable structures in numerical simulations (Wu 1988a; Steinolfson \\& Hundhausen 1990a,b). It has since been realised that the arguments against the existence of intermediate shocks are based on ideal MHD, which breaks down completely within the shock front. Numerical studies of time-dependent intermediate shocks in {\\it resistive} MHD demonstrated that the shocks could be formed `naturally' and were stable (Wu 1990). An examination of the structure of resistive, intermediate shocks in the weak shock limit (Kennel et al. 1989; Wu \\& Kennel 1992), shows that, as found for C-type shock waves, for given external parameters (i.e. pre-shock conditions and shock speed), a one-parameter family of intermediate shock structures exist connecting the jump conditions. This `internal' parameter provides the additional freedom for the shock to adjust its structure to external perturbations, a behaviour outside of the scope of ideal MHD in which the shock transition is represented as a discontinuity between the upstream and downstream states. As these conclusions depend on shock structure, in principle they depend on the model used for the magnetised medium. Resistive (Wu 1990) and hybrid (Wu \\& Hada 1991) models have been investigated so far. With the advent of multidimensional ambipolar diffusion codes, it has become possible to study these issues in weakly ionized media (see Smith \\& MacLow 1997). In this paper I constructed steady models of oblique C-type shocks in which the magnetic field within the shock front is not artificially confined to the plane containing the upstream and downstream magnetic field and the shock normal. Four fluids -- neutrals, ions, electrons or PAHs, and negatively charged grains -- were considered. The thermal pressure of the fluids and the inertia of the charged components were neglected. The effects of chemistry and changes in fractional ionization, and the charge residing on grains were ignored. The rate coefficients for grain-neutral and electron-neutral elastic scattering (which depend in principle on the neutral temperature and the grain-neutral drift speed, and on the electron temperature, respectively) were assumed constant with values appropriate to conditions within the shock front. These assumptions reduce the number of differential equations describing the shock structure to a pair for the components \\( B_x,\\,B_y \\) of the magnetic field that are perpendicular to the shock normal. Shock solutions can therefore be conveniently represented by trajectories in a two-dimensional phase space. Models were presented for weak shocks (\\( v_s/v_A = 1.5 \\)), and for strong shocks (\\( v_s/v_A =10 \\)). Variations in the grain size from \\( 0.1 \\mu \\) to \\( 0.4 \\mu \\), and the effect of an MRN grain-size distribution were considered. The results are summarized as follows: \\begin{enumerate} \\item The cold MHD jump conditions permit three stationary points in the phase space, corresponding to the upstream state and downstream states of the fast (for \\(v_s > v_A\\)) and intermediate shocks (\\(v_A\\,{\\rm cos}\\,\\theta < v_s < v_A {\\rm cot}\\,\\theta\\)). Valid shock solutions have phase-space trajectories linking the upstream state to one of the downstream states. A linear analysis, valid for \\emph{any} number of charged species, shows that the upstream state is a source, the fast downstream state is a saddle, and the intermediate downstream state is a sink. \\item A one-parameter family of intermediate shocks exists for each shock speed in the range \\(v_A\\,{\\rm cos}\\,\\theta < v_s < v_A {\\rm cot}\\,\\theta\\) where \\(\\theta\\) is the angle between the shock normal and the pre-shock magnetic field. These solutions correspond to those found by Pilipp \\& Hartquist (1994). The family contains members with either sense of rotation of \\( \\Bp \\) through the shock front, members with a rarefaction precursor in which \\( \\Bp \\) becomes small within the front before compression begins, and members that correspond to a fast shock followed downstream at a finite distance by an intermediate shock. \\item A unique fast shock exists for each shock speed \\(v_s > v_A\\), where \\(v_A\\) is the Alfv\\'en speed in the pre-shock gas. Pilipp \\& Hartquist (1994) were unable to find these solutions because integration of the equations for shock structure from the upstream state to the downstream state is unstable. Instead, integration must begin at the downstream state and run backwards through the shock front. \\item When all charged particles are well couwell coupled to the magnetic field, or when there is a symmetry between the poorly-coupled particles of either sign, the phase space is symmetric about the \\( B_x \\) axis, and the fast shock is coplanar, its trajectory running along the \\( B_x \\) axis. In conditions typical of molecular clouds, negatively-charged grains contribute significantly to the drag and are loosely coupled to the field. The net asymmetry in the coupling of negatively and positively charged particles to the magnetic field lines imposes a handedness on the shock structure and the phase space trajectories lose their symmetry on reflection in the \\( B_x \\) axis. In particular, the fast shock trajectory no longer runs along the axis -- the magnetic field and fluid velocities no longer lie in the \\(x\\)-\\(z\\) plane (containing the pre-shock field and shock normal) but have significant \\(y\\) components within the shock front. \\item For typical conditions in molecular clouds, grains dominate the frictional force and the magnetic field is well-tied to ions, electrons and PAHs. Under these conditions for given values of pre-shock magnetic field and gas density, the pre-shock grain Hall parameter determines the shock structure apart from an overall length scale for the shock thickness which is determined by the grain abundance. Thus under the approximations adopted in this paper, the shock structures are relatively insensitive to the details of the grain-size distribution and the exact composition of the charged species in the gas. \\item An MRN grain-size distribution can be approximately incorporated by calculating an effective abundance and an effective grain Hall parameter for single-size grains. The ratio of the effective abundance to the abundance of charged MRN grains, and the effective Hall parameter are determined by the Hall parameter of the smallest grains in the MRN distribution. \\item The degree of noncoplanarity is determined by the grain Hall parameter \\( \\beta_g \\), with significant effects being found for \\( |\\beta_g| \\la 1\\). When \\( |\\beta_g| \\la 0.3 \\) the upstream state becomes a spiral node and the shocks exhibit a precursor in which \\( \\Bp \\) may make several rotations. \\item Ions and electrons become passive in the sense that they may stream through the neutral gas with large transverse drift velocities (of order half the shock speed or more) without generating significant dissipation, as the shock thickness is determined by the collisions of neutrals with charged grains. The primary role of these charged species is to guarantee that the electric field in the shock front is very nearly orthogonal to the magnetic field. \\item Supressing the out-of-plane components of the drift velocities and magnetic field may significantly reduce the current and leads to a thicker shock structure (by, e.g. a factor of two) and a consequent decrease in the heating rate within the shock front. Models supressing the out-of-plane components of the magnetic field and velocity therefore underestimate the temperatures within the shock front, and significantly underestimate the magnitude of the grain drift speed through the neutrals. \\end{enumerate} This work was initiated at the University of Rochester. A. Perez-Miller is thanked for assistance with coding. This research was partially supported by NASA to the University of Rochester through grant NAGW-2444. The Special Research Centre for Theoretical Astrophysics is funded by the Australian Research Council under the Special Research Centre programme." + }, + "9802/astro-ph9802088_arXiv.txt": { + "abstract": "The distance to the centroid of the M31 globular cluster system is determined by fitting theoretical isochrones to the observed red-giant branches of fourteen globular clusters in M31. The mean true distance modulus of the M31 globular clusters is found to be $\\mu_0 = 24.47 \\pm 0.07$ mag. This is consistent with distance moduli for M31 that have been obtained using other distance indicators. ", + "introduction": "} Over the last ten years the distance to the Andromeda Galaxy ($=$~M31~$=$~NGC~224) has been measured using a variety of distance indicators. Pritchet~\\&~van~den~Bergh\\markcite{PV87}~(1987) used RR~Lyrae variables to derive a true distance modulus of $24.34 \\pm 0.15$. The observed brightnesses of red giants in the halo of M31 suggest that $\\mu_0 = 24.23 \\pm 0.15$ (Pritchet~\\&~van~den~Bergh\\markcite{PV88}~1988). Novae light curves give a distance modulus of $24.27 \\pm 0.20$ (Capaccioli~\\etal\\markcite{CD89}~1989) while Cepheid variables suggest that $\\mu_0 = 24.43 \\pm 0.06$ (Freedman~\\&~Madore\\markcite{FM90}~1990). Brewer~\\etal\\markcite{BR95}~(1995) found $\\mu_0 = 24.36 \\pm 0.03$ using carbon stars in the disk of M31 yet Ostriker~\\&~Gnedin\\markcite{OG97}~(1997) found $\\mu_0 = 24.03 \\pm 0.23$ from the peak of M31's globular cluster (GC) luminosity function. A weighted mean of these six values yields $\\langle\\mu_0\\rangle = 24.36 \\pm 0.03$ mag where the uncertainty is the standard error in the mean. This corresponds to a distance of $\\sim 745$ kpc. Recently Feast~\\&~Catchpole\\markcite{FC97}~(1997) used {\\sl Hipparcos\\/} data to calibrate the Cepheid period--luminosity (PL) relation and found a distance modulus of $\\mu_0 = 24.77 \\pm 0.11$ for M31. This corresponds to a distance of $\\sim 900$ kpc, an increase of $\\sim 20$\\% over the previously accepted value. However, Madore~\\&~Freedman\\markcite{MF97}~(1997) analysed the {\\sl Hipparcos\\/} data and found that the pre-{\\sl Hipparcos\\/} zero point for the Cepheid PL relation is good to within $\\pm 0.14$ mag ($\\sim 7$\\%). They also note that uncertainties due to reddening, metallicity effects on the PL relation, and statistical uncertainties in the data make up a significant component of the total uncertainty in the PL zero point. This study determines an estimate of the distance to M31 by finding the distance to the centroid of the M31 GC system. The shape of the red-giant branch (RGB) in a GC is strongly sensitive to the metallicity of the GC (see Da~Costa~\\&~Armandroff\\markcite{DC90}~1990 for a discussion of this effect in Galactic GCs). If the metallicity of the GC is determined in a manner that is {\\sl independent of the distance to the GC}, and the amount of reddening along the line-of-sight to the GC is known, then a reference isochrone of the correct metallicity can be shifted in distance until the best fit to the GC's RGB is obtained. ", + "conclusions": "} The distances to fourteen GCs in the M31 system have been estimated by comparing the observed RGB of each GC to a theoretical isochrone with the same iron abundance. The unweighted mean distance modulus of the GCs is $\\langle\\mu_0\\rangle = 24.47 \\pm 0.07$ mag. If only the ``best'' data set (the largest data set that used the same instrument and calibrations) is used a distance modulus of $\\langle\\mu_0\\rangle = 24.37 \\pm 0.06$ mag is obtained. These estimates of the distance to M31 are consistent with distances obtained using other techniques and do not support the need for a revision of the distance to M31. This method of determining the distance to M31 depends on accurate iron abundance estimates for the individual GCs and needs a sample of M31 GCs with high-precision photometry of their RGB stars. In order to improve the accuracy of this distance estimator deep ($V_{\\lim} \\gtrsim 26$) photometry will be needed for a much larger sample of M31 GCs. In addition, further work needs to be done to reduce the uncertainties in the spectroscopic iron abundance measurements of M31 GCs." + }, + "9802/astro-ph9802041_arXiv.txt": { + "abstract": "Using the latest compilation of cataclysmic variable orbital periods by Ritter \\&\\ Kolb we argue against Verbunt's conclusion that the period gap is not significant for nova-like variables. We also discuss the relation of the VY~Scl stars to the dwarf novae. ", + "introduction": "Two fundamental characteristics of a cataclysmic variable (CV) are the orbital period, $P$, typically between 80 m and 9 h, and the behaviour of the optical lightcurve, showing either the recurrent 2--5 mag outbursts of a dwarf-nova (DN), or the steady lightcurve of a nova-like variable (NL). DN are thought to be low accretion rate (\\mdot) systems with accretion discs cool enough to undergo hydrogen ionization instabilities, whereas NLs are always too hot for this to occur (see Warner 1995 for a comprehensive review of the whole field). The NLs predominantly have $P$\\,$>$\\,3 h, implying a high \\mdot\\ and thus high angular momentum loss from the binary in this range. Systems with $P$\\,$<$\\,2 h are predominantly DN, implying a lower angular momentum loss from gravitational radiation alone. Fewer systems occur in the `period gap' between 2 and 3 h, a possible consequence of the mechanism for additional braking switching off at \\sqig 3 h. However, Verbunt (1997) has concluded that the period gap is not significant for NLs, and thus queried whether the additional braking mechanism (usually suggested to be braking by the magnetic field of the secondary) is required. In this paper we re-examine the issue using Ritter \\&\\ Kolb's (1998) more recent list of CVs with known orbital period. Further, we address the nature of VY~Scl stars, which are currently poorly understood (e.g.~Livio \\&\\ Pringle 1994; Wu, Wickramasinghe \\&\\ Warner 1995; Verbunt 1997). ", + "conclusions": "This paper has been largely a defence of orthodoxy, showing that the period gap in cataclysmic variables is significant and that the nova-like variables show a cutoff at \\sqig 3 hr. Further, we've argued against Verbunt's (1997) classification of VY~Scl stars with the dwarf novae. Verbunt suggested that since VY~Scl variability is probably caused by changes in mass transfer from the secondary star, this could also have a major role in DN outbursts. We've shown that the differences are sufficient to require separate mechanisms, such as the disc instability for DN and irradiation-driven mass transfer cycles for VY~Scl stars." + }, + "9802/hep-ex9802007_arXiv.txt": { + "abstract": " ", + "introduction": "Double beta decay yields -- besides proton decay -- the most promising possibilities to probe beyond standard model physics beyond accelerator energy scales. Propagator physics has to replace direct observations. That this method is very effective, is obvious from important earlier research work and has been stressed, e.g. by \\cite{4}, etc.. Examples are the properties of $W$ and $Z$ bosons derived from neutral weak currents and $\\beta$--decay, and the top mass deduced from LEP electroweak radiative corrections. The potential of double beta decay includes information on the neutrino and sneutrino mass, SUSY models, compositeness, leptoquarks, right--handed $W$ bosons and others (see Table 1). The recent results of the Heidelberg--Moscow experiment, which will be reported here, have demonstrated that $0\\nu\\beta\\beta$ decay probes already now the TeV scale on which new physics should manifest itself according to present theoretical expectations. To give just one example, inverse double beta decay $e^-e^-\\rightarrow W^-W^-$ requires an energy of at least 4 TeV for observability, according to present constraints from double beta decay \\cite{14}. Similar energies are required to study, e.g. leptoquarks \\cite{3,5,hir96a,126,127,128,129}. To increase by a major step the present sensitivity for double beta decay and dark matter search, we present here for the first time a new project which would operate one ton of `naked' enriched {\\bf GE}rmanium detectors in liquid {\\bf NI}trogen as shielding in an {\\bf U}nderground {\\bf S}etup (GENIUS). It would improve the sensitivity from the present potential of at best $\\sim 0.1$ eV to neutrino masses down to 0.01 eV, a ten ton version even to 0.001 eV. The first version would allow to test a $\\nu_e \\rightarrow \\nu_{\\mu}$ explanation of the atmospheric neutrino problem, the second directly the large angle solution of the solar neutrino problem. The sensitivity for neutrino oscillation parameters would be larger than for all present accelerator neutrino oscillation experiments, or those planned for the future. GENIUS would further allow to test the recent hypothesis of a sterile neutrino and the underlying idea of a shadow world (see section 2). Both versions of GENIUS would definitely be a breakthrough into the multi-TeV range for many beyond standard models currently discussed in the literature, and the sensitivity would be comparable or even superior to LHC for various quantities such as right--handed W--bosons, R--parity violation, leptoquark or compositeness searches. Another issue of GENIUS is the search for Dark Matter in the universe. The full MSSM parameter space for predictions of neutralinos as cold dark matter could be covered already in a first step of the full experiment using only 100 kg of $^{76}$Ge or even natural Ge, making the experiment competitive to LHC in the search for supersymmetry. ", + "conclusions": "Double beta decay has a broad potential for providing important information on modern particle physics beyond present and future high energy accelerator energies which will be competitive for the next decade and more. This includes SUSY models, compositeness, left--right symmetric models, leptoquarks, and the neutrino and sneutrino mass. Based to a large extent on the theoretical work of the Heidelberg Double Beta group, results have been deduced from the HEIDELBERG--MOSCOW experiment for these topics and have been presented here. For the neutrino mass double beta decay now is particularly pushed into a key position by the recent possible indications of beyond standard model physics from the side of solar and atmospheric neutrinos, dark matter COBE results and others. New classes of GUTs basing on degenerate neutrino mass scenarios which could explain these observations, can be checked by double beta decay in near future. The HEIDELBERG--MOSCOW experiment has reached a leading position among present $\\beta\\beta$ experiments and as the first of them now yields results in the sub--eV range. We have presented a new idea and proposal of a future double beta experiment (GENIUS) with highly increased sensitivity based on use of 1 ton or more of enriched `naked' $^{76}$Ge detectors in liquid nitrogen. This new experiment would be a breakthrough into the multi-TeV range for many beyond standard models. The sensitivity for the neutrino mass would reach down to 0.01 or even 0.001 eV. The experiment would be competitive to LHC with respect to the mass of a right--handed W boson, in search for R--parity violation and others, and would improve the leptoquark and compositeness searches by considerable factors. It would probe the Majorana electron sneutrino mass more sensitive than NLC (Next Linear Collider). It would yield constraints on neutrino oscillation parameters far beyond all present terestrial $\\nu_e - \\nu_x$ neutrino oscillation experiments and could test directly the atmospheric neutrino problem and the large angle solution of the solar neutrino problem. GENIUS would cover the full SUSY parameter space for prediction of neutralinos as cold dark matter and compete in this way with LHC in the search for supersymmetry. Even if SUSY would be first observed by LHC, it would still be fascinating to verify the existence and properties of neutralino dark matter, which could be achieved by GENIUS. Concluding GENIUS has the ability to provide a major tool for future particle-- and astrophysics. Finally it may be stressed that the technology of producing and using enriched high purity germanium detectors, which have been produced for the first time for the Heidelberg--Moscow experiment, has found meanwhile applications also in pre-GENIUS dark matter search \\cite{101,102,Kla97e,Bau97} and in high--resolution $\\gamma$-ray astrophysics, using balloons and satellites \\cite{Kla91,81,109,100,111,Kla97b}." + }, + "9802/astro-ph9802292_arXiv.txt": { + "abstract": "The brown dwarf population in the Pleiades cluster has been probed in a deep 850 arcmin$^{2}$ $RIJK$ survey. The survey is complete to $I=21.4$ in 76\\,\\% of the area and to $I=20.2$ in the remaining 24\\,\\%. Photometry of 32 previously known members is presented together with 8 new candidates, four of which are below the brown dwarf limit. The faintest one is the lowest mass brown dwarf candidate found hitherto in the Pleiades ($I=20.55$, $0.04$ M$_{\\odot}$). The derived Pleiades luminosity function is compared to the most recent theoretical mass-luminosity relations and is consistent with a power-law index in the mass function between 0 and 1 to the limit of this survey. ", + "introduction": "The Pleiades cluster has been the major target of several recent surveys for brown dwarfs (BDs) (Hambly et al. 1993 (HHJ); Jameson \\& Skillen 1989 (JS); Schultz 1997; Stauffer et al. 1989, 1994a; Williams et al. 1996; Zapatero Osorio 1997; Zapatero Osorio et al. 1997a, b (ZRM, ZMR); Festin 1997; Cossburn et al. 1997). Its nearness ($116 \\pm 3$ pc, Mermilliod et al. \\cite*{mermilliod97}) and youth (120 Myr, Basri et al. \\cite*{basri96}) makes the rapidly cooling BDs still rather bright and easy to detect. The first bona fide Pleiades BD was reported by Rebolo et al. \\cite*{rebolo95}. This object, known as Teide1, passed the lithium test \\cite{rebolo96} and should thereby have a maximum mass of 0.06 M$_{\\odot}$. By now, on the order of 10 BDs have been confirmed in the Pleiades. The present observations were designed to probe low-mass BDs in the Pleiades and to provide accurate first-epoch data for future proper motion determination. This paper is an extension of the $IJK$ survey described in Festin \\cite*{festin97a}. \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{H0594_f1.ps}} \\caption{Histogram of seeing in the $I$ images} \\label{Figseeing} \\end{figure} ", + "conclusions": "We have performed a deep 850 arcmin$^{2}$ $RIJK$ survey in the central area of the Pleiades cluster. Photometry in $I$, $J$ or $K$ is presented for 32 previously known and 7 new likely members. Four of those are below the BD limit, and the faintest one would be the lowest mass BD found in the Pleiades so far, if its membership could be confirmed. The overall agreement in the photometry with other surveys is satisfactory. Teide1 is an exception. We find $I=19.26$, $\\sim$\\,0.5 mag fainter than ZRM. Based on RGO archive data and a roughly constant $J$ magnitude we conclude that the reason for this discrepancy most likely sits in the photometry of ZRM and not in intrinsic variability. A number of very red faint objects were found below the Pleiades sequence. Two of those were measured in $I$, $J$ and $K$ and fit in as GD165B-type objects, possible field BDs. After splitting PPL15 and two other probable binaries into components, the Pleiades LF was compared to model LFs derived from the most recent theoretical mass-luminosity relations. The observed LF supports an MF-index between 0 and 1. Thus, even if the MF seems to rise for low-mass BDs in the Pleiades, it is not steep enough to leave more than a few percent of the cluster's mass in BDs, which is consistent with dynamical findings \\cite{pinfield97b}." + }, + "9802/astro-ph9802198_arXiv.txt": { + "abstract": " ", + "introduction": "Most of the scientific community which has an interest in the physics of neutron stars believes that radio emission from Rotation Powered Pulsars (RPPs) has its origin in the relativistic outflow of electron-positron pairs along the polar magnetic field lines of a dipole magnetic field frozen into the rotating neutron star (e.g., Arons 1992, Meszaros 1992). The evidence for dipole magnetic fields in RPPs (and in any other neutron star) is indirect, coming primarily from the electromagnetic theory of RPP spindown. The observed increasing pulse periods are readily explained using standard theoretical moments of inertia plus order of magnitude estimates, derivable from dimensional analysis (Dyson 1971, Arons 1979, 1992), of rotational energy loss driven by relativistic electromagnetic spindown torques, \\begin{equation} \\dot{E}_R = k \\frac{\\mu^2 \\Omega_*^4}{c^3} = -I \\Omega_* \\dot{\\Omega}_*. \\label{rotloss} \\end{equation} Here $\\mu$ is the magnetic moment, $\\Omega_*$ is the stellar angular velocity with respect to inertial space far from the star, and $k$ is a function of any other parameters of significance, with magnitude on the order of unity. In the vacuum theory (Deutsch 1955), $k =(2/3) \\sin^2 i$, with $i$ the angle between the magnetic moment and the angular velocity. Theoretical work on the torques due to conduction currents steming back to Goldreich and Julian (1969), coupled to the observation that spindown rates appear to be independent of observationally estimated values of $i$ (Lyne and Manchester 1988), suggest that in reality $k$ does not substantially depend on $i$. In the subsequent discussion, I assume k = 4/9, the average of the vacuum value over the sphere. Application of (\\ref{rotloss}) to the observations of RPPs' periods ($P = 2 \\pi /\\Omega_*$ and period derivatives ($\\dot{P} = -2 \\pi \\dot{\\Omega}_* / \\Omega_*^2$) yields $\\mu \\sim 10^{30}$ cgs for ``normal'' RPPs, and $\\mu \\sim 10^{27}$ cgs for millisecond RPPs. These results are reasonably firm, the main uncertainty coming from the derived values of $\\mu$ being proportional to $k^{-1/2}$. The electromagnetic torque interpretation of pulsar spindown constrains only the exterior dipole moment of the magnetic field. However, not long ago Rankin (1990) presented strong evidence in favor of a low altitude ($r \\approx R_*$) dipole geometry for the site of the core component of pulsar radio emission. Arons (1993) gave evidence that spun up millisecond pulsars must have a substantially dipolar large scale field at low altitute. Electron-positron pair creation at low altitude above the polar caps has long been hypothesized to be an essential ingredient of pulsar radio emission, starting with Sturrock's (1971) pioneering work. If so, all observed pulsars must lie in the region of $P-\\dot{P}$ space where polar cap acceleration has sufficient vigor to lead to copious pair production. Yet, to date, all {\\it internally consistent} theories of polar cap pair creation have required hypothesizing a large scale ({\\it e.g.}, quadrupole) component of the magnetic field with strength comparable to that of the dipole (Ruderman and Sutherland 1975, Arons and Scharlemann 1979, Barnard and Arons 1982, Gurevich and Istomin 1985). These non-dipole components were invoked in order to increase the opacity of the magnetic field to pair creating gamma rays. Non-dipole low altitude fields can have magnetic radii of curvature on the order of $R_*$ or less, a factor of 50-100 smaller than the radii of curvature of star centered dipole field lines near the magnetic poles. The resulting increase of optical depth allowed the pair creation models to cover the whole $P, \\; \\dot{P}$ diagram. However, such strong magnetic anomalies contradict the evidence in favor of an apparently dipolar low altitude geometry; the alteration of the magnetic geometry also ruins the internal consistency of many models' electrodynamics. Both early (Sturrock 1971) and more recent work on polar cap electrodynamics and its implications for the occurrence of pair creation in $P, \\; \\dot{P}$ space either employ incomplete (e.g. Sturner {\\it et al.} 1995) or erroneous (Sturrock 1971, Mestel and Shibata 1994, Bjornsson 1996) theories of polar cap particle acceleration. Most of the internally consistent theories also violate other observational constraints, especially with regard to polar cap heating (Arons 1992), which creates pulsed thermal X-ray emission from hot spots in excess of what is seen (Becker and Tr\\\"{u}mper 1997, Pavlov and Zavlin 1997). While the Arons and Scharlemann (1979) model does not have this problem, in star centered dipole geometry it dramatically fails to account for pulsar emission over most of the $P-\\dot{P}$ diagram and predicts radio polarization variations in contradiction to the observations (Narayan and Vivekanand 1982). Here I describe a low altitude polar cap acceleration theory which successfully associates pulsar ``death'' with the cessation of pair creation in an {\\it offset} dipolar low altitude magnetic field. The basic acceleration physics is that of a space charge limited relativistic particle beam accelerated along the field lines by the starvation electric field, as in the Arons and Scharlemann theory, but with the additional effect of inertial frame dragging, first pointed out by Muslimov and Tsygan (1990, 1992) and by Beskin (1990). This effect causes the accelerating electric field to be about an order of magnitude larger than that calculated by Arons and Scharlemann for pulsars near the death line, which substantially improves the size of the region in $P, \\; \\dot{P}$ space in which polar cap pair creation occurs, but still does not allow the theory to fully account for the observed pulsar distribution, in star centered dipole geometry. If the dipole's center is offset from the stellar center along a vector parallel to the dipole moment itself, an offset which automatically preserves the symmetries built into the highly successful Radhakrishnan and Cooke (1969) model of polarization swings, the magnetic field at one pole becomes substantially stronger than it would be if the same magnetic dipole were star centered. If the offset is substantial (as much as 80\\% of the stellar radius turns out to be required), all pulsars can be accommodated within a single pair creation theory. The location of an individual pulsar's pair creation death depends on the magnitude of the offset, thus yielding a ``death valley'' (Chen and Ruderman 1993) for the whole pulsar population. ", + "conclusions": "I have shown that polar pair creation based on acceleration of a steadily flowing, space charged limited non-neutral beam in a locally dipolar magnetic geometry at low altitude is consistent with pulsar radio emission throughout the $P - \\dot{P}$ diagram, provided 1) the effect of dragging of inertial frames is included in estimates of the starvation electric field; 2) the dipole center is strongly offset from the stellar center, perhaps as much as $0.7-0.8 R_*$; and 3) inverse Compton emission of thermal photons from a neutron star cooling slower than exponentially at ages in excess of $10^6$ years plays an important role in the emission of magnetically convertible gamma rays. The development of new diagnostics of the low altitude magneic field, and gamma ray observations sensitive to low altitude emission, will eventually provide tests of these ideas." + }, + "9802/astro-ph9802151_arXiv.txt": { + "abstract": "Near-infrared spectroscopy ($\\lambda_{\\rm rest} \\sim${\\ts}3700--6800{\\ts}\\AA) of eight high redshift powerful radio galaxies (HzPRGs) at $z = 2.2-2.6$ is presented. Strong forbidden lines and H$\\alpha$ emission were detected in all sources; the data show evidence that the emission lines of the HzPRGs may contribute a substantial fraction ($\\sim$ {\\ts}25--98{\\ts}\\%) of their total observed $H$- and/or $K$-band light. Diagnostic emission-line ratios -- [O III]{\\ts}$\\lambda$5007{\\ts}/{\\ts}H$\\beta$ vs. [S II]{\\ts}$\\lambda \\lambda$6716, 6731{\\ts}/{\\ts}H$\\alpha$ -- for three of the eight HzPRGs are consistent with the presence of a Seyfert{\\ts}2 nucleus; the [O III]{\\ts}$\\lambda$5007{\\ts}/{\\ts}H$\\beta$ and [S II]{\\ts}$\\lambda \\lambda$6716, 6731{\\ts}/{\\ts}H$\\alpha$ ratios and/or limits of the remaining five galaxies are inconclusive. Furthermore, all six of the galaxies for which both $H$- and $K$-band spectra were obtained have observed [O III]{\\ts}$\\lambda$5007{\\ts}/{\\ts}(H$\\alpha$+[N II] $\\lambda \\lambda$6548, 6583) ratios consistent with Seyfert{\\ts}2 ionization. Much of the inability to detect the weaker emission lines of [S II]{\\ts}$\\lambda \\lambda$6716, 6731 in three of the galaxies and H$\\beta$ in any of the galaxies may be due to moderate amounts of dust: for the two sources with previously measured Ly$\\alpha$ fluxes, the observed Ly$\\alpha${\\ts}/{\\ts}H$\\alpha$ ratios are $\\sim${\\ts}1.5, much less than the value of 16 expected for gas in a dust-free medium photoionized by a hard, nonthermal continuum. If such a discrepancy is due solely to dust, this ratio translates into $A_V \\sim${\\ts}0.5--1.0 mag (depending on which extinction curve -- Milky Way, SMC, LMC -- is used) at the rest-frame optical wavelengths of the galaxies, and a corresponding factor of $\\sim${\\ts}1.6--2.5 reduction in optical flux. None of the eight HzPRGs at $z = 2.2-2.6$ HzPRGs have broad ($\\Delta v_{\\rm FWHM} > 1500$ km s$^{-1}$) emission-line cores, and it is not clear whether any have broad emission-line wings. However, the near-infrared spectrum of 3C{\\ts}22, a $z = 0.937$ radio galaxy with 1{\\ts}$\\mu$m luminosity comparable to that of the radio galaxies at $z=2.2-2.6$ and a radio luminosity only 3--5 times less, shows direct evidence for broad H$\\alpha$ emission wings. Such a feature is indicative of the presence of a partially obscured Seyfert{\\ts}1 nucleus. Given that 3C{\\ts}22 is at $\\sim${\\ts}1/3 the luminosity distance of the sample of HzPRGs at $z = 2.2-2.6$, a thorough search for such a faint feature in the more distant galaxies may require 8--10 meter-class telescopes. These new data, along with recent UV-to-optical polarimetry showing evidence of high polarization in many HzPRGs, provide evidence that many HzPRGs are predominantly ionized by an active nucleus, and that a significant fraction of their SED may be due to non-thermal emission from an active galactic nucleus. ", + "introduction": "High redshift powerful radio galaxies (HzPRGs: $P_{\\rm 408 MHz} > 10^{27} h^{-2}${\\ts}W Hz$^{-1}$, $z \\gtrsim${\\ts}1)\\footnote{Throughout this paper, $H_o = 100 h${\\ts}km s$^{-1}$ Mpc$^{-1}$ and $q_o = 0.5$ are assumed.} are potentially an important probe of galaxy evolution at early epochs. The recent detections of strong emission lines from HzPRGs at observed near-infrared wavelengths (e.g., Rawlings, Eales, \\& Lacy 1991; McCarthy, Elston, \\& Eisenhardt 1992; Eales \\& Rawlings 1993, 1996) has provided a new diagnostic tool with which to study these objects, although the nature of line emission in HzPRGs is currently a subject of great debate. Part of the difficulty in interpreting line emission in HzPRGs has been that the use of standard optical diagnostic emission-line ratios (e.g. Baldwin, Phillips, \\& Terlevich 1981; Veilleux \\& Osterbrock 1987), commonly used to differentiate between thermal ionization and low or high-excitation ionization from a non-thermal source, has until recently been problematic because rest-frame optical emission from galaxies at $z >${\\ts}1 is redshifted to near-infrared wavelengths at current epochs. A new generation of infrared spectrographs has now made it possible to begin systematic studies of the dominant ionization mechanisms in HzPRGs via their rest-frame optical emission-line spectra. In particular, the new {\\it K}-band spectrograph (KSPEC: Hodapp et al. 1994) on the University of Hawaii (UH) 2.2{\\ts}m telescope on Mauna Kea provides the unique capability of simultaneous coverage of the $\\sim${\\ts}1.0--2.4{\\ts}\\micron$~~$ wavelength band at typical spectral resolution $\\lambda${\\ts}/{\\ts}$\\Delta \\lambda$ {\\ts}$\\sim${\\ts}700, while the infrared spectrograph (CGS4) on the 3.8{\\ts}m United Kingdom Infrared Telescope (UKIRT) on Mauna Kea can provide coverage of either the full {\\it J}, {\\it H}, or {\\it K}-band near-infrared windows at a spectral resolution of $\\lambda${\\ts}/{\\ts}$ \\Delta \\lambda${\\ts}$\\sim${\\ts}860. A nearly complete sample of HzPRGs in the redshift range $z=2.2-2.6$ has been selected. An attempt has been made, insomuch as possible, to obtain near-infrared spectra with similar sensitivity and coverage in rest-frame wavelength. The choice of the redshift range and wavelength coverage was designed to maximize the number of rest-frame optical diagnostic lines observed. As far as the author is aware, this is the first such study that obtains both {\\it H}- and {\\it K}-band spectra for a fairly complete sample of HzPRGs. The outline of this paper is as follows: Sample selection and observing procedures are discussed in \\S{\\ts}2 and \\S{\\ts}3 respectively. Data reduction methods are summarized in \\S{\\ts}4. The discussion in \\S{\\ts}5 focuses on the observed emission-line spectra and calculated extinction in the HzPRGs, as well as on how the properties of HzPRGs compare with the properties of local radio galaxies, and optically selected Seyferts, LINERs, and H{\\ts}II region-like galaxies. ", + "conclusions": "\\subsection{Emission-Line Diagnostics} Figure 3 shows the log ([O III]{\\ts}$\\lambda$5007{\\ts}/{\\ts}H$\\beta$) vs. log ([S II]{\\ts}$\\lambda$6724{\\ts}/{\\ts}H$\\alpha$) diagram commonly used to distinguish between galaxies with Seyfert, LINER, and H{\\ts}II-region emission-line spectra. Emission from [S II]{\\ts}$\\lambda$6724 can emanate from ionized hydrogen regions, as well as from semi-ionized regions where collisional ionization is significant. Enhancement of forbidden lines such as [O I]{\\ts}$\\lambda$6300, [N II]{\\ts}$\\lambda \\lambda$6548, 6583, and [S II]{\\ts}$\\lambda$6724 occurs in AGNs because, unlike H{\\ts}II regions, they have extended partially-ionized zones created by an excess of X-rays (the absorption cross sections of neutral hydrogen, helium and all ions are small for X-rays, thus X-rays tend to escape the ionized region before interacting: Veilleux \\& Osterbrock 1987). [O III]{\\ts}$\\lambda$5007 is a high ionization line photoionized by UV photons, and thus tends to be strong in Seyfert galaxies. Instead of attempting to determine the relative contributions of H$\\alpha$ and the [N II]{\\ts}$\\lambda \\lambda$6548, 6583 doublet to the blended H$\\alpha$+[N II]{\\ts}$\\lambda \\lambda$6548, 6583 complex, the smallest ratio of [N II]{\\ts}$\\lambda$6583{\\ts}/{\\ts}H$\\alpha$ (= 0.19) observed for a sample of luminous infrared galaxies (Kim et al. 1995) has been adopted and plotted (large circles with embedded numerals) for all of the HzPRGs in the sample\\footnote{ It is unlikely, given the widths of the [O III]{\\ts}$\\lambda$5007 lines, that the H$\\alpha$+[N II]{\\ts}$\\lambda \\lambda$6548, 6583 complexes are pure H$\\alpha$. However, if the complexes were purely H$\\alpha$, their observed breadth ($\\Delta v_{\\rm FWHM} \\gtrsim$ 1000 km s$^{-1}$) would be evidence against H{\\ts}II region-like emission.}. Choosing such a ratio tends to bias the data towards the H{\\ts}II-region portion of the diagram. To illustrate the extent to which the data are affected by this ratio, the values that result from assuming [N II]{\\ts}$\\lambda$6583{\\ts}/{\\ts}H$\\alpha$ = 1.0, the value adopted by Eales \\& Rawlings (1993) based on the predictions of photoionization models and the observed value for low-redshift powerful radio galaxies (LzPRGs), and [N II]{\\ts}$\\lambda$6583{\\ts}/{\\ts}H$\\alpha$ = 3.74, the largest value observed by Kim et al. (1995), have also been plotted as small, filled-in circles. Note that in all of the $H$-band spectra except for the spectrum of 4C{\\ts}48.48, the estimated H$\\beta$ upper limits (3{\\ts}$\\sigma$) provide weak constraints on the lower limit of [O III]{\\ts}$\\lambda$5007{\\ts}/{\\ts}H$\\beta$. Thus, in all sources except 4C{\\ts}48.48 and 4C{\\ts}40.36, the [O III]{\\ts}$\\lambda$5007{\\ts}/{\\ts}H$\\beta$ ratio is determined assuming H$\\alpha${\\ts}/{\\ts}H$\\beta$ $\\ge${\\ts}3, and thus [O III]{\\ts}$\\lambda$5007{\\ts}/{\\ts}H$\\beta$ $\\ge${\\ts}3{\\ts}$\\times${\\ts}[O III]{\\ts}$\\lambda$5007{\\ts}/{\\ts}H$\\alpha$ (the upper limit for [O III]{\\ts}$\\lambda$5007{\\ts}/{\\ts}H$\\beta$ in 4C{\\ts}40.36 is taken from Iwamuro et al. 1996). Three of the eight HzPRGs plotted in Figure 3 clearly fall in the Seyfert region of the plot. The [O III]{\\ts}$\\lambda$5007{\\ts}/{\\ts}H$\\beta$ and [S II]{\\ts}$\\lambda$6724{\\ts}/{\\ts}H$\\alpha$ ratios and/or limits for TX{\\ts}0200+015, TX{\\ts}0828+193, 3C{\\ts}257, MG{\\ts}1744+18, and 4C{\\ts}23.56 are inconclusive. Additional information about the ionization mechanism is obtained by examining the [O III]{\\ts}$\\lambda$5007{\\ts}/{\\ts}(H$\\alpha$+[N II] $\\lambda \\lambda$6548, 6583) ratios. Table 3 lists the average value for this ratio as a function of emission-line classification for galaxies in the {\\it IRAS} Bright Galaxy Sample (BGS) and a sample of ``warm'' {\\it IRAS} galaxies (Kim et al. 1995; Veilleux et al. 1995), as well as for a sample of LzPRGs with Seyfert{\\ts}2 emission-line spectra (i.e., Cygnus A: Osterbrock \\& Miller 1975; PKS{\\ts}1345+12: Grandi 1977; 3C98, 3C192, 3C327: Costero \\& Osterbrock 1977). Despite the obvious scatter due presumably to extinction and possible metallicity effects, the ratio provides a notable separation between Seyfert galaxies and those of the H{\\ts}II and LINER class. A comparison of these low-redshift, active galaxies with the six HzPRGs at $z = 2.2-2.6$ for which both $H$- and $K$-band spectra have been obtained show all six HzPRGs to have [O III]{\\ts}$\\lambda$5007{\\ts}/{\\ts}(H$\\alpha$+[N II] $\\lambda \\lambda$6548, 6583) ratios consistent with Seyfert{\\ts}2 galaxies\\footnote{The [O III]{\\ts}$\\lambda$5007 measurement of 4C{\\ts}40.36 by Iwamuro et al. (1996), obtained with a $PA = 0$ and a 1.5\\arcsec$\\times$60\\arcsec$\\,$ aperture, is included because their $H$-band image shows most of the flux to be within 1.5\\arcsec$\\,$ of the center of the galaxy. The [O III]{\\ts}$\\lambda$5007 flux is $3.86\\times10^{-18}$ W m$^{-2}$, and thus the observed [O III]{\\ts}$\\lambda$5007{\\ts}/{\\ts}H$\\alpha$+[N II] $\\lambda \\lambda$6548, 6583 ratio for 4C{\\ts}40.36 is 1.1.}. The average value of the [O III]{\\ts}$\\lambda$5007{\\ts}/{\\ts}(H$\\alpha$+[N II] $\\lambda \\lambda$6548, 6583) ratio for these six HzPRGs is also listed in Table 3. Thus, the dominant source of ionization for these galaxies appears to be no different from that observed in most low-redshift, narrow-line radio galaxies. Because the H$\\alpha$ and [N II]{\\ts}$\\lambda \\lambda$6548, 6583 lines are blended, a direct comparison of the low ionization-to-H$\\alpha$ line ratio of these high-redshift galaxies to their low-redshift counterparts cannot be made. An alternative is to compare the [S II]{\\ts}$\\lambda$6724{\\ts}/{\\ts}(H$\\alpha$+[N II] $\\lambda \\lambda$6548, 6583) ratio. For the 5 HzPRGs with [S II]{\\ts}$\\lambda$6724 detections, [S II]{\\ts}$\\lambda$6724{\\ts}/{\\ts}(H$\\alpha$+[N II] $\\lambda \\lambda$6548, 6583) {\\ts}$= 0.25 \\pm 0.07$, similar to the mean value of $0.27 \\pm 0.05$ observed for a sample of LzPRGs with Seyfert{\\ts}2 emission-line spectra. Changes in properties such as the elemental abundances of the host galaxy, the shape of the ionizing spectrum, and the electron density of the gas being ionized will act to vary this ratio (note that extinction has little effect on this ratio because these emission lines have similar wavelengths). Table 3 contains a summary of average [S II]{\\ts}$\\lambda$6724{\\ts}/{\\ts}(H$\\alpha$+[N II] $\\lambda \\lambda$6548, 6583), H$\\alpha${\\ts}/{\\ts}[S II]{\\ts}$\\lambda$6724, H$\\alpha${\\ts}/{\\ts}[N II] $\\lambda \\lambda$6548, 6583, and [N II]{\\ts}$\\lambda \\lambda$6548, 6583{\\ts}/{\\ts}[S II]{\\ts}$\\lambda$6724 ratios as a function of emission-line classification. Essentially, [S II]{\\ts}$\\lambda$6724{\\ts}/{\\ts}(H$\\alpha$+[N II] $\\lambda \\lambda$6548, 6583) is a function of H$\\alpha${\\ts}/{\\ts}[S II]{\\ts}$\\lambda$6724 and [N II]{\\ts}$\\lambda \\lambda$6548, 6583{\\ts}/{\\ts}[S II]{\\ts}$\\lambda$6724, where H$\\alpha${\\ts}/{\\ts}[S II]{\\ts}$\\lambda$6724 appears to account for most of the variation in the average. On average, [S II]{\\ts}$\\lambda$6724{\\ts}/{\\ts}(H$\\alpha$+[N II] $\\lambda \\lambda$6548, 6583) is larger for LINERs and Seyfert 2 galaxies than in H{\\ts}II region-like galaxies (see Table 3), in part because the former have more extended semi-ionized regions (i.e., H$\\alpha${\\ts}/{\\ts}[S II]{\\ts}$\\lambda$6724 decreases as the extent of the semi-ionized region increases). Variations in the metal abundance in the galaxies in the sample undoubtedly contribute to the scatter. For example, a substantial decrease in the metal abundance would cause a substantial decrease in [S II]{\\ts}$\\lambda$6724{\\ts}/{\\ts}(H$\\alpha$+[N II] $\\lambda \\lambda$6548, 6583) (i.e., H$\\alpha${\\ts}/{\\ts}[S II]{\\ts}$\\lambda$6724 increases as the metal abundance decreases), and vice versa. The observation that LzPRGs and HzPRGs have similar [S II]{\\ts}$\\lambda$6724{\\ts}/{\\ts}(H$\\alpha$+ [N II]{\\ts}$\\lambda \\lambda$6548, 6583) ratios may indicate that, even though HzPRGs are at substantially higher lookback times and possess higher radio luminosities, the properties of the ionizing source and the ionized gas are similar. It is worthwhile to consider whether the gas in the HzPRGs could be ionized by supernova remnants (SNRs). Indeed, there are a few SNRs that exhibit emission-line ratios similar to Seyfert{\\ts}2 galaxies. A conservative estimate of the SNR rate required to explain the observed emission lines can be determined by considering the extreme SNR in NGC{\\ts}6946, which has log([O III]{\\ts}$\\lambda$5007{\\ts}/{\\ts}H$\\beta$) = 0.85 and log([S II]{\\ts}$\\lambda$6724{\\ts}/{\\ts}H$\\alpha$) $= -0.06$ (Blair \\& Fesen 1994). The flux in the [O III]{\\ts}$\\lambda$5007 line is $2.5\\times10^{-17}$ Watts m$^{-2}$. At a distance corresponding to $z = 2.4$, the SNR would have a flux of $\\sim 2.5\\times10^{-17}$(5.1 Mpc / $9330h^{-1}$ Mpc)$^2$ = $7.5\\times10^{-24}h^2$ Watts m$^{-2}$. Thus, to produce the average [O III]{\\ts}$\\lambda$5007 flux (i.e. $3\\times10^{-18}$ Watts m$^{-2}$) observed in the HzPRGs in the sample would require $4\\times10^5h^{-2}$ SNRs. Adopting a conservative estimate of 5000{\\ts}yr for the SNR lifetime, HzPRGs would have to produce SNRs at an implausible rate of $80h^{-2}${\\ts}yr$^{-1}$ to maintain their [O III]{\\ts}$\\lambda$5007 flux. \\subsection{$L_{\\rm [O III]{\\ts}\\lambda \\lambda4959, 5007}$ vs. $P_{\\rm 151 MHz}$} Previous authors have used the [O III]{\\ts}$\\lambda \\lambda$4959, 5007 luminosity versus the observed 151 MHz radio power for HzPRGs in an attempt to infer a causal relationship between the ionization source for the gas and the source of the radio emission. Rawlings et al. (1989) have compiled data primarily for LzPRGs, and more recently, Eales \\& Rawlings (1993, 1996) have added data for HzPRGs. Figure 4a is a plot adapted from Eales \\& Rawlings (1996) of the [O III]{\\ts}$\\lambda \\lambda$4959, 5007 luminosity versus the observed 151 MHz radio power for an unbiased 3C sample of FR{\\ts}II radio galaxies with $z <${\\ts}0.5 and a collection of HzPRGs. Data for the HzPRGs TX{\\ts}0200+015, B3{\\ts}0731+438, TX{\\ts}0828+193, 4C{\\ts}48.48, and 4C{\\ts}23.56 have also been plotted. A simple interpretation of Figure 4a would be that there appears to be a correlation between the [O III] $\\lambda \\lambda$ 4959, 5007 emission-line luminosity and the radio power (such a correlation would appear to be even tighter perhaps for the HzPRGs than for the LzPRGs), and that a common excitation source (e.g. the central AGN) is responsible for both. However, the real answer is clearly not so simple. The apparent strong correlation between the [O III] $\\lambda \\lambda$ 4959, 5007 emission-line luminosity and the radio power, $P_{\\rm 151 MHz}$, in Figure 4a is almost entirely an artifact of distance, as evidence by the plot of the ratio $L_{\\rm [O III]{\\ts}\\lambda \\lambda4959, 5007} / P_{\\rm 151 MHz}$ vs. $P_{\\rm 151 MHz}$ (Figure 4b) which shows no correlation in the sample as a whole. However, note the apparent strong correlation exhibited in Figure 4b by the HzPRGs (i.e., all sources with $P_{\\rm 151 MHz} > 10^{27}$ W Hz$^{-1}$ sr$^{-1}$ and $L_{\\rm [O III]{\\ts}\\lambda \\lambda4959, 5007} / P_{\\rm 151 MHz} > 10^8$). Indeed, for the most powerful radio galaxies at $z >${\\ts}0.5 the morphology of the emission-line gas is aligned with the radio axis (McCarthy et al. 1987), which suggests that a correlation between the two quantities exists. The extraordinary strength of the emission lines in HzPRGs is also evident from the data in Table 4: of the TX and 4C galaxies in the sample for which high signal-to-noise ratio $H$- or $K$-band magnitudes were made available from Keck imaging (L. Armus, private communication), the emission lines appear to contribute anywhere between 25--98{\\ts}\\% of the broad-band near-infrared light. These percentages are substantially higher than those observed in the most powerful radio galaxies in the local Universe, but are consistent with other HzPRGs with published near-infrared spectra (McCarthy et al. 1992; Eales \\& Rawlings 1993, 1996). Such high emission-line luminosities illustrate the danger of interpreting the magnitudes and morphologies of high radio-power galaxies as being purely stellar in origin (see Eales \\& Rawlings 1993, 1996 for discussions on this issue ). Evidence for possible AGN contamination of the broad-band morphologies of high radio-power galaxies has been demonstrated by Dunlop \\& Peacock (1993), who show that the rest-frame 1.1 $\\mu$m flux of a sample of 3CR galaxies tend to be more extended and aligned with the radio axis than a sample of lower radio-power Parkes galaxies in the same redshift range ($z \\sim 1$). Such extended and aligned morphologies are more pronounced at rest-frame UV wavelengths and are undoubtedly connected with the energetics of the central engine, and/or due to an optically-thick, circumnuclear dust/gas torus extinguishing light perpendicular to the radio axis. Rest-frame UV spectropolarimetry (Dey \\& Spinrad 1995; Cimatti et al. 1996; Manzini \\& di Serego Alighieri 1996 and references therein) and optical imaging polarimetry (Knopp \\& Chambers 1997; Knopp 1997), show further evidence of AGN contamination of broad-band light in many HzPRGs.\\footnote{There is one example of an HzPRG, 4C 41.17 at $z=3.8$, that does not appear to have polarized UV light (Dey et al. 1997).} Indeed, four of the galaxies discussed in this paper (TX{\\ts}0200+015, B3{\\ts}0731+438, TX{\\ts}0828+193, and 4C{\\ts}23.56) show rest-frame optical polarizations of 10--45\\% (Knopp \\& Chambers 1997; Knopp 1997), indicating that a significant fraction of the rest-frame UV-to-optical continuum emission from many HzPRGs may be scattered/reprocessed AGN light. \\subsection{TX{\\ts}0828+193SW: A Component Dominated by Continuum-Emission} In some HzPRGs, there appear to be ``components'' that are genuinely dominated by continuum emission at observed optical and near-infrared wavelengths. Figure 1 shows two extractions of the southwestern ``component'' of TX{\\ts}0828+193. Unlike the northeastern component, TX{\\ts}0828+193SW has no notable strong emission lines, but has comparatively strong continuum emission. Spectroscopy of the radio galaxy at wavelengths near 4000 \\AA$~$ (observer-frame) reveals strong Ly$\\alpha$, C{\\ts}IV, He{\\ts}II, and C{\\ts}III] emission lines in the NE component, but no emission lines in the SW component (van Ojik 1995; R\\\"{o}ttgering et al. 1997). Radio, $R$-band imaging data, and near-infrared imaging data taken of TX{\\ts}0828+193 show the radio core to be coincident with the NE component and the SW component to lie $\\sim$ 5\\arcsec~ away from the radio core, but still along the radio axis of the galaxy (R\\\"{o}ttgering et al. 1995; Knopp \\& Chambers 1997; see Figure 5 in the Appendix). This raises the possibility that either TX{\\ts}0828+193SW is a cloud being illuminated by continuum emission emanating from TX{\\ts}0828+193NE, or that a buried AGN residing in TX{\\ts}0828+193SW is ionizing gas in TX{\\ts}0828+193NE. The two best examples of ionization of off-nuclear knots in low-redshift radio galaxies are Coma A (van Breugel et al. 1985) and PKS 2152-69 (Tadhunter et al. 1987). In the case of Coma A, the nuclear region has strong continuum emission and relatively weak lines, whereas the off-nuclear knot has relatively strong line emission. In the case of PKS 2152-69, both the nucleus and the off-nuclear knot have notable continuum emission, but the line emission from the off-nuclear knot, especially [O III]$\\lambda$ 5007, is stronger. Such data would imply that the AGN is in TX{\\ts}0828+193SW and is ionizing the northeastern component. Recent multi-wavelength, broad-band polarimetry measurements of TX{\\ts}0828+193 show evidence for TX{\\ts}0828+193NE having polarization consistent with scattering due to dust (Knopp \\& Chambers 1997). Because TX{\\ts}0828+193SW has no emission-line or absorption features from which its redshift can be determined, there exists the possibility that it is simply an unrelated object along the line of sight. Though the chance of such a superposition is low, there is precedent for concern, most notably the foreground star coincident with 3C368 (Hammer, Le F\\`{e}vre, \\& Proust 1991). In the recently discovered HzPRG MG{\\ts}1019+0535, a double-component morphology is observed, one component of which shows no sign of strong line emission (Dey, Spinrad, \\& Dickinson 1995). The authors argue strongly in favor of the lineless component being a foreground object, but unlike TX{\\ts}0828+193, the two `components' of MG{\\ts}1019+0535 are orthogonal to the radio axis. Given the geometry of TX{\\ts}0828+193, the polarimetry measurements, and that such double-component structures aligned with the radio axis are observed in other radio galaxies (e.g. MRC{\\ts}0406-24: Eales \\& Rawlings 1993), TX{\\ts}0828+193SW is most likely at the redshift of the radio source. \\subsection{Dust} The evidence to date that HzPRGs as a class contain substantial amounts of dust is mixed, but compelling - while polarimetry observations show evidence for scattering by dust in several HzPRGs (Dey \\& Spinrad 1995; Cimatti et al. 1996; Manzini \\& di Serego Alighieri 1996 and references therein; Knopp \\& Chambers 1997; Knopp 1997), far-infrared/submillimeter bolometry (Golombek, Miley, \\& Neugebauer 1988; Evans et al. 1996; Dunlop et al. 1994; Chini \\& Kreugel 1994; Ivison 1995; Hughes, Dunlop, \\& Rawlings 1997) and CO spectroscopy (Evans et al. 1996; Downes et al. 1996; van Ojik et al. 1997a; Scoville et al. 1997) surveys of HzPRG show evidence for substantial amounts of dust in only a few sources. Regardless of whether or not HzPRGs have comparable or more dust as that inferred by the SEDs and molecular gas masses of many low-redshift radio galaxies (e.g, Golombek, Miley, \\& Neugebauer 1988; Mirabel, Sanders \\& Kaz\\`{e}s 1989; Knapp \\& Patten 1991; Impey \\& Gregorini 1993; Mazzarella et al. 1993; Evans 1996), only moderate amounts are required to notably decrease the observed flux of the optical (rest-frame) emission lines and continuum. Specifically, the extinction in the HzPRGs may be substantial enough such that the weaker lines of H$\\beta$, [O I]{\\ts}$\\lambda$6300, and [S II]{\\ts}$\\lambda$6724, and possible features such as emission-line wings (see \\S 5.5), fall below the detection threshold of the spectra in Figure 1. Estimates of extinction in HzPRGs in the sample were made by comparing the observed line ratios of hydrogen recombination lines with their intrinsic ratio (i.e., ratios in a dust-free environment). Intrinsic ratios for low-density gas photoionized by a thermal continuum source are H$\\alpha${\\ts}/{\\ts}H$\\beta$ = 2.85 and Ly$\\alpha${\\ts}/{\\ts}H$\\alpha$ = 8.10 (Osterbrock 1989). However, because the evidence presented here and elsewhere (see McCarthy 1993) implies that gas in HzPRGs is heated by a hard nonthermal continuum, enhanced Ly$\\alpha$ and H$\\alpha$ emission in predominantly neutral regions heated by X-rays must be taken into account (note that collisions resulting in the emission of Ly$\\alpha$ and H$\\alpha$ are comparatively more frequent than those resulting in H$\\beta$). Thus, the intrinsic line ratios applicable to HzPRGs are H$\\alpha${\\ts}/{\\ts}H$\\beta$ = 3.1 and Ly$\\alpha${\\ts}/{\\ts}H$\\alpha$ = 16 (Ferland \\& Osterbrock 1985; Osterbrock 1989). The H$\\beta$ emission line was not detected in any of the HzPRG spectra shown in Figure 1. Thus, the Ly$\\alpha${\\ts}/{\\ts}H$\\alpha$ ratio must be used to calculate the extinction for individual sources. This ratio can be determined for four of eight sources; in the case of B3{\\ts}0731+438 ($\\sim 2.6$) and 3C{\\ts}257 ($\\sim 0.1$) by using data from Eales \\& Rawlings (1993)\\footnote {McCarthy et al. (1992) have also computed Ly$\\alpha${\\ts}/{\\ts}H$\\alpha$ for B3{\\ts}0731+438, as well as for the HzPRG MRC 0406-24. The discrepancy between the ratios determined by them and those determined by Eales \\& Rawlings (1993) arise from the fact that the former assume H$\\alpha =$ H$\\alpha$+[N II] and the latter assume H$\\alpha = \\case{1}{2}$ (H$\\alpha$+[N II]).}, and in the case of TX{\\ts}0200+015 ($\\sim 1.7$) and TX{\\ts}0828+193 ($\\sim 1.4$) by using data from Table 2 and R\\\"{o}ttgering et al. (1997). The average observed Ly$\\alpha${\\ts}/{\\ts}H$\\alpha$ value determined for a larger sample of lower-redshift 3CR radio galaxies is $\\sim${\\ts}1.5 (McCarthy 1988), and the observed Ly$\\alpha${\\ts}/{\\ts}H$\\alpha$ ratio determined for two other $z>2$ HzPRGs (Eales \\& Rawlings 1993) are $\\sim 1.4$ (MRC{\\ts}0406-24) and $\\sim 1.4$ (53W002). All of the observed ratios are much less than the value of 16 expected for gas in a dust-free medium photoionized by a hard, nonthermal continuum. The $E(B-V)$ values determined for the $z>2$ HzPRGs B3 0731+438, 3C{\\ts}257, MRC 0406-24, and 53W002 using data from Eales \\& Rawlings (1993) are in the range 0.11--0.76\\footnote{This is the range of $E(B-V)$ for these sources assuming the extinction curves for the Milky Way, the Large Magellanic Cloud, and the Small Magellanic Cloud, and assuming an [N II]{\\ts}$\\lambda$6583{\\ts}/{\\ts}H$\\alpha$ ratio of 1.0 and 0.19. See the rest of the paragraph for a detailed explanation.}. Using the H$\\alpha$ measurements in Table 2, in combination with the R\\\"{o}ttgering et al. (1997) Ly$\\alpha$ measurements for TX 0200+015 and TX 0828+193, the $E(B-V)$ of these two galaxies was determined using the following procedure: Because the spectral resolution is sufficient to split the H$\\alpha$+[N II]{\\ts}$\\lambda$6583 complex, the values of H$\\alpha$ have been determined using the three ratios of [N II]{\\ts}$\\lambda$6583{\\ts}/{\\ts}H$\\alpha$ plotted on Figure 3 (i.e., 0.19, 1.0, 3.74). The observed Ly$\\alpha${\\ts}/{\\ts}H$\\alpha$ line ratios were first corrected for Galactic extinction, then the extinction in the HzPRGs was determined assuming the extinction curves for the Milky Way (Savage \\& Mathis 1979), LMC (Nandy et al. 1981), and SMC (Prevot et al. 1984); the corresponding values of $E(B-V)$ are listed in Table 5. Given the redshift of the H$\\alpha$+[N II]{\\ts}$\\lambda \\lambda$6548, 6583 complex and shape of the line emission, it seems quite unlikely that the [N II]{\\ts}$\\lambda$6583{\\ts}/{\\ts}H$\\alpha$ ratio is as high as 3.74. Thus, using the values from Table 5 and excluding values calculated assuming [N II]{\\ts}$\\lambda$6583{\\ts}/{\\ts}H$\\alpha$ = 3.74, feasible values of $E(B-V)$ range from 0.14 to 0.39. Such $E(B-V)$ values translate into $A_V \\sim 0.5-1.0${\\ts}mag (depending on which extinction curve -- Milky Way, SMC, LMC -- is adopted) at the observed infrared wavelengths (i.e., rest-frame optical wavelengths of HzPRGs at $z=2.2-2.6$), and a corresponding factor of $\\sim${\\ts}2 reduction in the emission-line and continuum flux. Although extinction is commonly determined using the luminous Ly$\\alpha$ and H$\\alpha$ emission lines (as done above), there are three concerns in using Ly$\\alpha$ for such a measurement. These concerns and their relative applicability to the analysis of HzPRGs are summarized here. First, if there is a sufficient density of atomic hydrogen, the number of Ly$\\alpha$ photons along the line of sight can be greatly diminished by resonant scattering or associated absorption. One argument against this is that the measured Ly$\\alpha$ emission-line widths are broad enough ($\\gtrsim$ few hundred km s$^{-1}$) such that most of the photons are in the wings of the line (McCarthy 1996). This can be understood if there exists a large velocity gradient across the gas that is being ionized - Ly$\\alpha$ radiation emitted from gas at a given radius from the AGN passes through the gas farther out which is traveling at slower velocities because, to this gas, the Ly$\\alpha$ photons do not appear to be resonant photons. However, recent medium resolution spectroscopy of 15 HzPRGs (van Ojik et al. 1997b), of which TX 0200+015 and TX 0828+193 are included, show evidence in favor of Ly$\\alpha$ absorption. If the intrinsic Ly$\\alpha$ profiles they fit to the observations are correct, the Ly$\\alpha$ emission in these two galaxies may be diminished by up to 50\\%, making $E(B-V)$ lower by 0.10, 0.07, and 0.04 for models assuming a Milky Way, a LMC, and a SMC extinction curve, respectively. Second, if the gas is dense enough, Ly$\\alpha$ photons can be destroyed by collisional de-excitation of the 2$^2P$ state. This is because the Ly$\\alpha$ photons must random walk out of the gas, and thus the lifetime of the 2$^2P$ state is lengthened by the number of steps a Ly$\\alpha$ photon must take to escape the gas. However, the densities required for the collisional de-excitation timescale to be shorter than the radiative de-excitation lifetime of the 2$^2P$ state are in excess of 10$^{10}$ cm$^{-3}$, a density only believed to be found in the densest areas of the broad-line region of AGNs (Osterbrock 1989). Third, if the alignment of the rest-frame UV morphology and the radio axis, as observed in many HzPRGs at $z >${\\ts}0.5, is the result of external illumination of gas and dust clouds by radiation from a central engine, these clouds most likely reflect the majority of the Ly$\\alpha$ photons (e.g., McCarthy 1996). Given such a geometry, the Ly$\\alpha$ luminosity would provide little information on the amount of dust present. However, such a process would also tend to raise, not lower, the Ly$\\alpha$/H$\\alpha$ ratio from its intrinsic value, which is the opposite of what is observed. \\subsection{Broad Lines} None of the eight HzPRGs at $z = 2.2-2.6$ were found to have broad emission-line cores (see Figure 1), nor do they appear to have broad emission-line wings. However, it is entirely possible that the cosmological distances of these galaxies may simply limit the ability to detect such a feature with 2--4 meter-class telescopes. The emission-line spectrum of the lower redshift HzPRG 3C{\\ts}22 suggests that this may indeed be the case. Figure 2 shows that the H$\\alpha$+[N II]{\\ts}$\\lambda \\lambda$6548, 6583 complex for 3C{\\ts}22 has a line core width, $\\Delta v_{\\rm FWHM} \\sim$ 1700 km s$^{-1}$, consistent with the higher redshift HzPRGs observed (see Table 2). This emission emanates from extended, low density gas some distance from the nucleus of the galaxy. However, note the very broad H$\\alpha$ wings ($\\Delta v_{\\rm FWZI} \\sim$ 7600 km s$^{-1}$) in 3C{\\ts}22; such a feature is also visible in additional data (not shown here) from three KSPEC observing periods between 1994, July--September, and has also been reported by Economou et al. (1995) and Rawlings et al. (1995). Such a feature is indicative of the presence of a partially obscured Seyfert{\\ts}1 nucleus, and is consistent with the hypothesis that FR{\\ts}II radio galaxies are quasars where the broad-line active nucleus is mostly obscured from the line of sight. The signal-to-noise ratio of the spectra of the higher redshift galaxies shown in Figure 1 is simply insufficient to rule out the presence of a similarly broad component in the $z = 2.2-2.6$ objects. It is tempting to speculate that broad line wings, such as those observed in 3C{\\ts}22, may be present in the more distant HzPRGs. However, current efforts using 2--4 meter-class telescopes suggests that larger aperture (i.e., 8--10 meter-class telescopes) will be required to detect broad line wings for HzPRGs at $z > 2$." + }, + "9802/astro-ph9802221_arXiv.txt": { + "abstract": "s{ In inflationary cosmology, the particles constituting the Universe are created after inflation in the process of reheating due to the interaction with the oscillating inflaton field. We briefly review the basics of the slow reheating, and the stage of fast preheating, when the particles are created explosively in the regime of parametric resonance. The non-perturbative, out-of-equilibrium character of the parametric resonance changes many features of reheating. For these proceedings, we will highlight a few aspects of preheating: the structural dependence of the parametric resonance on the inflationary model, including $V(\\phi)={m^2 \\over 2}\\phi^2$, ${\\lambda\\over 4}\\phi^4$, $1-\\cos {\\phi \\over f}$; ``rescattering'' of created particles; and phase transitions after inflation.} ", + "introduction": " ", + "conclusions": "" + }, + "9802/astro-ph9802003_arXiv.txt": { + "abstract": "The recent discovery that the UV dust extinction in starburst galaxies is similar to that found in the Small Magellanic Cloud (SMC) motivated us to re-investigate the ultraviolet (UV) extinction found in the SMC. We have been able to improve significantly on previous studies by carefully choosing pairs of well matched reddened and unreddened stars. In addition, we benefited from the improved S/N of the NEWSIPS {\\it IUE} data and the larger sample of SMC stars now available. Searching the {\\it IUE} Final Archive, we found only four suitable early-type stars that were significantly reddened and had well matched comparison stars. The extinction for three of these stars is remarkably similar. The curves are roughly linear with $\\lambda^{-1}$ and have no measurable 2175~\\AA\\ bump. The fourth star has an extinction curve with a significant 2175~\\AA\\ bump and weaker far-UV extinction. The dust along all four sightlines is thought to be local to the SMC. There is no significant Galactic foreground component. The first three stars lie in the SMC Bar and the line-of-sight for each of them passes through regions of recent star formation. The fourth star belongs to the SMC Wing and its line-of-sight passes though a much more quiescent region. Thus, the behavior of the dust extinction in the SMC supports a dependence of dust properties on star formation activity. However, other environmental factors (such as galactic metallicity) must also be important. Dust in the 30 Dor region of the LMC, where much more active star formation is present, does not share the extreme extinction properties seen in SMC dust. ", + "introduction": "The interstellar dust in the Small Magellanic Cloud (SMC) has gained new importance with the discovery that its ultraviolet (UV) extinction is uniquely similar to extinction found in starburst galaxies (\\cite{cal94}; Gordon, Calzetti, \\& Witt 1997). Dust in the Galaxy and the Large Magellanic Cloud does not show starburst galaxy-like extinction. Starburst galaxies are the only type of galaxy that has been detected from low to high ($z > 2.5$) redshifts (\\cite{kin93}; \\cite{ste96}; \\cite{low97}; \\cite{tra97}; \\cite{fra97}). These galaxies serve as an excellent probe of galaxy evolution through the study of their star formation and metal enrichment rates at different redshifts. The derivation of these rates for high redshift galaxies is sensitive to the adopted UV dust extinction (\\cite{pet97}; \\cite{mad97}). Thus, understanding the physical processes responsible for producing the starburst-like dust seen in the SMC is important to the modeling of starburst galaxies. In addition, the study of the UV extinction curve in Local Group galaxies promises to help in determining the sizes, shapes, and materials which make up dust grains. Ultraviolet extinction curves have been determined in the Milky Way, the Large Magellanic Cloud (LMC), the SMC, and (tentatively) M31. The extinction curves in these four galaxies paint a complex picture of the environmental dependence of dust properties. In the Milky Way, Cardelli, Clayton, \\& Mathis (1989, hereafter \\cite{car89}) found that the infrared to UV extinction curves can be described fairly well by a relationship which depends on only one parameter, $R_V = A_V/E(B-V)$ which is a measure of the overall dust grain size. There are significant deviations from the \\cite{car89} relation in different Galactic environments (\\cite{mat92}). In the LMC, the UV extinction curves show a distinctly different behavior between the 30~Dor region (a mini-starburst [\\cite{wal91}]) and the rest of the LMC (\\cite{cla85}; \\cite{fit85}, 1986). The 2175~\\AA\\ bump is weaker and the far-UV rise is stronger in the 30~Dor region than in the rest of the LMC which have strengths similar to the average Galactic extinction curve. In the SMC, the average extinction curve is characterized by a roughly linear rise (versus $\\lambda^{-1}$) increasing toward shorter wavelengths without a 2175~\\AA\\ bump (\\cite{pre84}; \\cite{tho88}). Yet, there is one sightline which has an extinction curve with a significant 2175~\\AA\\ bump (\\cite{leq82}). In M31, the extinction curve is consistent with that of the average Galactic extinction within the associated uncertainties, although the 2175~\\AA\\ bump may be weak (\\cite{bia96}). The complex behavior in these four galaxies implies that the physical properties of dust grains may be dependent on a multitude of environmental parameters. Two of these are metallicity and star formation activity both of which may affect the overall composition and size distribution of dust grains (\\cite{cla85}; \\cite{fit85}; \\cite{car89}; \\cite{gor97}). ", + "conclusions": "} Any discussion of the behavior of the dust in the SMC based on only four extinction curves is obviously severely hampered by the small sample size. Yet, interesting trends can be seen even in this small sample. The four extinction curves are plotted together in Figure~\\ref{fig_all_elv_ebv} and their FM fit parameters are tabulated in Table~\\ref{table_FM_param}. The extinction curves for AZV~18, 214, \\& 398 are very similar and are all roughly linear with $\\lambda^{-1}$. The extinction curve for AZV~456 is much different as it has a significant 2175~\\AA\\ bump and a weaker far-UV extinction. The similarity of the extinction curves for AZV~18 \\& 214 (lightly reddened) to that of AZV~398 (more reddened) gives confidence that the extinction curves for AZV~18 \\& 214 are real. \\begin{figure}[tbp] \\begin{center} \\plotone{all_elv_ebv.eps} \\caption{The four extinction curves for AZV~18, 214, 398, \\& 456 are plotted. \\label{fig_all_elv_ebv}} \\end{center} \\end{figure} \\begin{deluxetable}{lcccccc} \\footnotesize \\tablewidth{0pt} \\tablecaption{FM fit parameters \\label{table_FM_param}} \\tablehead{\\colhead{AZV} & \\colhead{$c_1$} & \\colhead{$c_2$} & \\colhead{$c_3$} & \\colhead{$x_o$} & \\colhead{$\\gamma$} & \\colhead{$c_4$}} \\startdata 18 & $-5.68 \\pm 0.28$ & $2.53 \\pm 0.05$ & $0.76 \\pm 1.38$ & $4.56 \\pm 0.23$ & $1.68 \\pm 1.52$ & $0.60 \\pm 0.26$ \\nl 214 & $-3.72 \\pm 0.29$ & $2.03 \\pm 0.05$ & $0.09 \\pm 3.03$ & $4.98 \\pm 1.97$ & $5.12 \\pm 19.24$ & $-0.20 \\pm 0.26$ \\nl 398 & $-5.16 \\pm 0.29$ & $2.27 \\pm 0.05$ & $0.14 \\pm 11.29$ & $4.41 \\pm 17.06$ & $-6.48 \\pm 211.22$ & $0.80 \\pm 0.28$ \\nl 456 & $-0.96 \\pm 0.09$ & $1.18 \\pm 0.02$ & $2.57 \\pm 0.22$ & $4.71 \\pm 0.01$ & $1.00 \\pm 0.05$ & $0.10 \\pm 0.15$ \\nl \\enddata \\end{deluxetable} Since the extinction curve for AZV~456 looks more like Milky Way dust than SMC dust, one wonders if a large fraction of this extinction is due to foreground Galactic dust. Velocity resolved H I measurements toward AZV~456 show that 90\\% of the \\ion{H}{1} along the line-of-sight is located in the SMC (\\cite{leq82}; \\cite{mcg81}, 1982). Also, the star AZV~454, which is $1\\farcm 5$ from AZV~456, shows very little evidence of reddening. These observations strongly imply that most of the interstellar medium along the line of sight toward AZV~456 lies in the SMC. In the Milky Way, most of the differences between extinction curves can be explained by the variation in the single parameter $R_V = A(V)/E(B-V)$. This parameter measures the average dust grain size with $R_V$ increasing with increasing average grain size (\\cite{car89}). Following Bouchet et al.\\ (1985), the $R_V$ values for the four SMC extinction curves were calculated from \\begin{equation} \\label{eq_rv} R_V = 1.10 \\frac{\\Delta (V - K)}{\\Delta (B - V)}. \\end{equation} The intrinsic colors of the reddened stars were assumed to be either the colors of their respective comparison stars or the intrinsic colors of Galactic stars of the same spectral types as tabulated in Johnson (1966) for $(B-V)$, and Koornneef (1983), for $(V-K)$. The values of $R_V$ calculated both ways are tabulated in Table~\\ref{table_rv_values}. The two $R_V$ values are equivalent within the uncertainties and the adopted values are contained in the last column of Table~\\ref{table_rv_values}. The four extinction curves have roughly similar values of $R_V$. If the SMC dust followed a CCM-like relationship then AZV~456 would have the largest $R_V$. This does not seem to be the case. \\begin{deluxetable}{lccc} \\tablewidth{0pt} \\tablecaption{$R_V$ Values \\label{table_rv_values}} \\tablehead{\\multicolumn{1}{c}{star} & \\colhead{comparison} & \\colhead{Galactic} & \\colhead{adopted} } \\startdata AZV 18 & $3.60 \\pm 0.73$ & $2.78 \\pm 0.34$ & $3.60 \\pm 0.73$ \\nl AZV 214 & $2.75 \\pm 0.55$ & $2.36 \\pm 0.21$ & $2.75 \\pm 0.55$ \\nl AZV 398 & $2.87 \\pm 0.40$ & $3.05 \\pm 0.17$ & $2.87 \\pm 0.40$ \\nl AZV 456 & \\nodata & $2.66 \\pm 0.16$ & $2.66 \\pm 0.16$ \\nl \\enddata \\end{deluxetable} In order to investigate the dependence of the extinction in the SMC on environment, we plotted the positions of the four stars on an H$\\alpha$ image of the SMC (Figure~\\ref{fig_smc_Halpha}). The H$\\alpha$ intensities trace star formation activity. The lines-of-sight toward all four stars are associated with the known \\ion{H}{2} regions (\\cite{dav76}; \\cite{cap96}). The three stars with roughly linear extinction curves (AZV~18, 214, \\& 398) are located in regions of high H$\\alpha$ intensities (SMC Bar). The one star (AZV~456) with a more Galactic type extinction curve is also located in an \\ion{H}{2} region but one with much weaker star formation (SMC Wing). The \\ion{H}{2} region associated with the line-of-sight to AZV~214 is noteworthy because it is associated with the cluster NGC~346 which is the most massive star formation region in the SMC. NGC~346 contains 33 known O type stars (\\cite{mas89}). The H$\\alpha$ flux from this cluster is 10\\% of that seen from 30~Dor in the LMC (\\cite{mas89}). The H$\\alpha$ fluxes from the \\ion{H}{2} regions along the lines-of-sight toward AZV~18, 398, \\& 456 are are 1, 1, and 0.1\\%, respectively, of that seen for 30~Dor (\\cite{ken86}; \\cite{cap96}). The dust along the AZV~456 sightline has likely been exposed to a less harsh environment than the other three sightlines. \\begin{figure}[tbp] \\begin{center} \\plotone{smc_Halpha.eps} \\caption{The positions of the four reddened stars are plotted on a H$\\alpha$ image of the SMC (\\cite{bot97}). This image is displayed in the North Celestial Pole projection (\\cite{sta97}). The image was provided by G.\\ Bothun and L.\\ Staveley-Smith. \\label{fig_smc_Halpha}} \\end{center} \\end{figure} It is known that processing of Galactic dust near regions of active star formation results in changes in the UV extinction curve (\\cite{mat92}). A similar behavior is seen in the 30 Dor region in the LMC (\\cite{fit86}). It is clear from this work that the harsh radiation and shock environment associated with star formation regions is modifying the dust in the SMC. In all three galaxies, star formation modifies nearby dust by altering the 2175~\\AA\\ bump and increasing the strength of the far-UV extinction. Yet, the extent of the modification does {\\em not} seem to be well correlated with the level of star formation activity. Figure~\\ref{fig_all_ext} plots the average extinction curves for the LMC (30~Dor and rest of LMC, [\\cite{fit86}]), the SMC (Bar and Wing), and the Milky Way ($R_V = 3.1$). The extinction curve for $\\theta^1$~Ori~D, a star in the Orion \\ion{H}{2} region, is also plotted. The average extinction curve for the SMC Bar is similar to the previous SMC average extinction curve (\\cite{pre84}). From this figure, it is obvious that the extinction in the SMC Bar is unlike that which has been found anywhere except starburst galaxies (\\cite{gor97}). On the other hand, the extinction in the SMC Wing is similar to that found in the LMC (excluding the 30~Dor region) and the Milky Way open cluster Trumpler~37 (not shown, \\cite{cla87}; \\cite{fit90}). The large difference between the extinction curve in the Orion \\ion{H}{2} region and the SMC Bar implies that an Orion-like \\ion{H}{2} region is not enough to produce starburst-like dust. Something more is needed. \\begin{figure}[tbp] \\begin{center} \\plotone{all_ext.eps} \\caption{The extinction curves for the SMC, LMC, and Milky Way are plotted. The curves plotted are those calculated from the FM fits except for the Milky Way which was calculated from the CCM relationship for an $R_V = 3.1$. The extinction curve for the SMC Bar is the $\\Delta (B-V)$ weighted average of the curves for AZV~18, 214, \\& 398. The SMC Wing and $\\theta^1$~Ori~D extinction curves have been multiplied by 0.83 and 1.3, respectively, to allow easier comparison to the other four curves. \\label{fig_all_ext}} \\end{center} \\end{figure} This raises the question: Why are the extinction properties of the dust in the SMC Bar so much more extreme than that found in the 30~Dor region of the LMC? The largest star formation region in the SMC (NGC~346) has only 10\\% the activity of 30~Dor as measured by H$\\alpha$ (\\cite{cap96}). The extinction curve for AZV~456 shows that dust, very similar to Galactic and LMC dust, exists in the SMC. Therefore, there must be some significant environmental difference between the LMC 30~Dor and SMC Bar regions which affects the extent to which star formation activity can modify dust. One known difference between the LMC and SMC is metallicity. The metallicities of the LMC and SMC are 0.2 and 0.6 dex lower, respectively, than the that of the local Galactic ISM (the ISM is 0.1 dex lower than solar). However, the relative abundances of the elements in the LMC and SMC are similar to that found in the local interstellar medium (\\cite{rus92}). Metallicity is correlated with the dust-to-gas ratio in galaxies (\\cite{iss90}). So, the amount of dust in the SMC is significantly lower than that found in the LMC. This could affect the ability of dust grains in the SMC to shield themselves from the radiation and shocks present near star formation regions. This contrasts with the finding that starburst galaxies all possess dust with an extinction curve like that found in the SMC Bar even though they have metallicities between 0.1 and 2.0 solar (\\cite{gor97}). It is possible that the much higher level of star formation present in starburst galaxies (10$\\times$ that of 30~Dor) overwhelms other environmental factors and always produces SMC Bar-like dust. The starburst galaxies were UV selected biasing the sample toward intrinsically bright, nearby starbursts with at least one lightly reddened starburst region. Thus, the combination of a small column of dust and the more intense star formation possibly accounts for the presence of SMC-like dust in all starburst galaxies studied in Gordon et al.\\ (1997). } $\\bullet$ We have greatly improved the UV extinction curves for the SMC through improvements in the S/N of the IUE spectra and careful choices of reddened and comparison star pairs. $\\bullet$ Four reddened SMC stars possess the Fitzpatrick criteria needed for accurate extinction calculations. Three of the four stars possess a roughly linear (with $\\lambda^{-1}$) extinction curve. The lines-of-sight toward these three stars (AZV~18, 214, \\& 398) pass through active star formation regions. The fourth star (AZV~456) has an extinction curve with a 2175~\\AA\\ bump and a weaker far-UV rise. Its sightline also passes through a star formation region, but one which is much less active. $\\bullet$ Processing of dust near regions of star formation results in variations in UV extinction. However, there is no simple correlation between the strength of the variations and the amount of star formation activity. Other parameters such as galaxy metallicity must also play an important role. $\\bullet$ As this work is based on only four sightlines, more observations of reddened stars in the SMC are needed for to confirm these results. Of special need are observations outside the SMC Bar to confirm the extinction curve of AZV~456." + }, + "9802/astro-ph9802235_arXiv.txt": { + "abstract": "Type I planetary nebulae (PNe) are defined as those with high He and N abundances (Peimbert $\\&$ Torres-Peimbert 1983). These objects present in general bipolar geometries and have high stellar temperatures (Corradi $\\&$ Schwarz 1995, Torres-Peimbert $\\&$ Peimbert 1997). In this paper we analyse the empirical methods for abundance determination in order to check if the He and N overabundances in Type I PNe are a consequence of a geometrical effect, due to the bipolarity, or the ionization stratification, due to the stellar temperature. For this, we obtain simulated spherically symmetrical as well as bipolar nebulae, using a 3D photoionization code. From the projected emission line intensities for: a) the whole nebula; b) for a slit crossing the nebula; as well as c) for different positions in the nebula, we applied the formulae used in the literature to obtain empirical abundances. These empirical abundances are then compared with the adopted ones. We show that empirical abundances depend on the particular line of sight covered by the observation and can simulate an overabundance and/or the presence of abundance gradients of He and N in planetary nebulae with high stellar temperature. The geometrical effects are also discussed. Systematic errors in abundance determinations by empirical methods are higher for the N/H ratio than for N/O. Thus, it seems better to use the N/O value when discussing N rich objects. ", + "introduction": "The knowledge of chemical abundances in Planetary Nebulae (PNe) is important since they can be used as constraints for stellar evolutionary models, as well as for studies of the chemical evolution of galaxies. In fact, PNe abundances reflect the products of stellar evolution, as well as the enrichment of the interstellar medium. Chemical abundances in PNe are usually obtained using empirical methods (as for example in Peimbert $\\&$ Torres-Peimbert 1987, Kingsburgh $\\&$ Barlow 1994). Empirical methods are commonly used since they can easily be applied to a great number of objects, from a small number of bright emission lines. Photoionization models also provide the gas chemical composition when applied to specific nebulae (for example, Harrington et al. 1982; Clegg et al. 1987; Gruenwald, Viegas $\\&$ Brogui\\`ere 1997). Following their abundance, PN are classified in types (Peimbert 1978; Peimbert $\\&$ Torres-Peimbert 1983; Fa\\'undez-Abans $\\&$ Maciel 1987). Type I PNe are those with high He and N abundances, e.g. He/H $\\geq$ 0.125 and log(N/O) $>$ -0.30 (Peimbert $\\&$ Torres-Peimbert 1983). In a study of morphological and physical properties of PNe, Corradi $\\&$ Schwarz (1995) state that bipolar nebulae have the hottest stars among PNe, and, except for two objects, all bipolar PNe for which chemical data are available are Type I. Another common characteristic of bipolar nebulae is that the [NII] line intensity is stronger than H$\\alpha$ all over the nebula, apart from the very central region (Corradi $\\&$ Schwarz 1993b). This characteristic is usually explained by a N overabundance or by shock heating (Corradi $\\&$ Schwarz 1993a; Corradi $\\&$ Schwarz 1993c). Predictions of stellar evolutionary models and the high values obtained for He and N abundances in Type I/bipolar planetaries have suggested that their central stars are more massive than other PN nuclei (Peimbert $\\&$ Serrano 1980; Calvet $\\&$ Peimbert 1983; Torres-Peimbert 1984; Corradi $\\&$ Schwarz 1995). However, new evolutionary models fail to explain the high abundances of He and N in Type I PNe (Marigo, Bressan $\\&$ Chiosi 1996), although models including the ``hot botton burning'' effect and accounting for the chemical evolution on the AGB may lead to higher N/O abundances (van den Hoek $\\&$ Groenewegen 1997). From the characteristics of Type I/bipolar planetary nebulae, it can be noted that objects with the highest He and N empirical abundances are those presenting a bipolar geometry and having the highest stellar temperatures. In fact, objects with higher stellar temperatures (T$_*$) tend to present higher He and N empirical abundances (see Fig. 1). In order to homogenize the data in the figure, only objects with abundances obtained from line intensities integrated in a slit across the nebulae are plotted. As it will be shown in the following sections, empirical abundance determinations depend on the line of sight covered by the observation. Empirical abundances for the objects plotted in Fig. 1 are from de Freitas Pacheco et al. (1992), Kingsburgh $\\&$ Barlow (1994), and Costa et al. (1996). The stellar temperature corresponds to the Zanstra temperature of HeII. Values for T$_*$ are taken from the literature (Kaler 1983; Preite-Martinez $\\&$ Pottasch 1983; Pottasch 1984; Shaw $\\&$ Kaler 1985; Gathier $\\&$ Pottasch 1988; Gathier $\\&$ Pottasch 1989; Gleizes, Acker, $\\&$ Stenholm 1989; Kaler $\\&$ Jacoby 1989; Shaw $\\&$ Kaler 1989; Kaler, Shaw, $\\&$ Kwitter 1990; Stanghellini, Corradi, $\\&$ Schwarz 1993). In this paper we analyze the effect of the geometry and of the stellar temperature on the empirical abundance determinations, in order to verify if the high abundances of He and N obtained for Type I PNe are real. In fact, several studies show that the stellar temperature and/or the geometry can affect the abundances obtained from empirical methods. As discussed by Peimbert and collaborators (Peimbert 1967; Peimbert $\\&$ Costero 1969; Peimbert, Luridiana, $\\&$ Torres-Peimbert 1995), the gas temperature obtained from observed forbidden line ratios and used in abundance determinations must be corrected by temperature fluctuations inside the nebulae. In fact, for typical conditions of PNe, temperature fluctuations can be important when the ionizing central star is hotter than 10$^5$ K (Gruenwald $\\&$ Viegas 1995). Furthermore, a study on spherically symmetrical homogeneous HII regions show that temperature fluctuations obtained from projected data vary with the line of sight to the nebula; the variation depending on the characteristics of the star and the nebula (Gruenwald $\\&$ Viegas 1992). With a three-dimensional (3D) self-consistent photoioinization code, a model for the bipolar PN IC 4406 has been obtained, showing that its abundance does not fulfill the criteria for Type I, as this nebula is usually classified (Gruenwald et al. 1997). Such an analysis shows that specific and more realistic models which take into account the geometry of the object can lead to abundance results that differ from those obtained with empirical methods. In order to check for the effects of geometry and stellar temperature on the empirical method to obtain PNe chemical abundances, we apply our self-consistent 3D photoionization code to planetary nebulae. For a wide range of nebular and stellar characteristics, we obtain the physical conditions in each point of the nebula, and calculate the resulting line intensity ratios at different lines of sight to the nebula. From the calculated line intensity ratios we determine the abundances applying the empirical methods used in the literature. These abundances are then compared to those assumed in the models. The models are described in \\S 2. The results are presented in \\S 3 and the conclusions are outlined in \\S 4. ", + "conclusions": "The results of the preceding section show that empirical methods for abundance determination can simulate an overabundance and/or the presence of abundance gradients of He and N in planetary nebulae with high T$_*$, even when applied to spherically symmetrical and homogeneous objects. This is also true for N/O. Recall that Type I PNe, defined as those nebulae which present an overabundance of He and N, have the hottest stars among PNe, and are generally bipolar. As shown in \\S 3, both characteristics can induce an apparent overabundance when empirical methods are used. Following our results (Figs. 3a and 3b), systematic errors in the abundance determinations by empirical methods are smaller for N/O. Thus, it is better to use the N/O ratio instead of N/H when discussing N rich planetaries. For central stellar temperatures in the range 10$^5$ to 3$\\times$10$^5$ K the N/O ratio can be overestimated by a factor of 1.5. On the other hand, the average N/O value for Type I PNe is about 2.8 higher than Type II PNe (Peimbert 1990). So, we expect that, even with a better N/O determination, Type I PNe may still be N rich compared to Type II PNe. Inference about abundances on planetary nebulae are important tests to stellar evolutionary models. Several arguments indicate that Type I PNe come from more massive progenitors (Torres-Peimbert $\\&$ Peimbert 1997). In particular, high masses for the central stars of Type I PNe are generally invoked to explain their high He and N abundances since the higher N/O ratio, the higher the original mass of the central star. Some theoretical evolutionary models fail to explain these overabundances (Marigo et al. 1996). However, models accounting for the chemical evolution on AGB and including the \"hot bottom burning\" effect are consistent with PNe abundances (van den Hoek $\\&$ Groenewegen 1997). Our results show that He and N overabundances in Type I PNe may be lower than indicated by the empirical methods. The correction of the empirical abundances by those systematic errors change the constraints to the evolutionary models and may lead to different conclusions on the importance of the different processes. Furthermore, from the discussion in the preceding section, empirical methods applied to PNe with high stellar temperatures can mimic abundance gradients. Thus abundance gradients and their relation to multiple ejection phenomena must also be reviewed. Let us recall that observed line intensities are integrated on the line of sight, and different lines of sight cross different fractions of regions differing by their ionization state. Since lines originate in different regions, the observed image of a nebula depends on choice of the emission line. As shown in the preceding section, abundances obtained from observed line intensities depend on the line of sight. Therefore, empirical abundances obtained from line intensities averaged over different points in a nebula, or conclusions based on average abundances for a given nebula, may be incorrect. The effects shown in the present paper must be taken as a warning when discussing chemical enrichment and/or abundance gradients in a given nebula. Any conclusion about abundances, in particular in nebulae with a high star temperature, even with simple geometries, must rely on detailed photoionization models which properly account for the integration of emission line intensities on a given line of sight. Classical unidimensional photoionization models are suitable for spherically symmetrical objects, whereas a 3D photoionization code (Gruenwald at. al 1997) must be used for other geometries." + }, + "9802/astro-ph9802145_arXiv.txt": { + "abstract": " ", + "introduction": "In a hybrid thermal-nonthermal steady-state plasma, thermal and nonthermal electrons coexist. We Assume that the electron distribution is isotropic and can be described by \\begin{equation} f( \\gamma ) =\\left \\{ \\begin{array}{ll} f_{th}( \\gamma )=A \\gamma \\sqrt{\\gamma^{2}-1} e^{-\\frac{( \\gamma-1 )}{T_{e}}}& \\gamma \\leq \\gamma_{th} \\\\ f_{nth}( \\gamma )= B \\gamma^ {-p} & \\gamma \\geq \\gamma_{th} \\end{array} \\right . \\end{equation} where $\\gamma$ is the Lorentz factor, $T_e$ is the temperature of thermal electrons in unit of $mc^2$, $f(\\gamma)$ is normalized to 1, $A$ and $B$ are determined by the normalization and the continuity condition at $\\gamma_{th}$. This hybrid distribution may represent the particle distribution of many astronomical sources. From many RS CVn binary systems, good correlations of radio emission with x-ray and UV emission are observed (Drake, Simon \\& Linsky 1989, 1992). Spectra show that the x-ray emission is due to thermal bremsstrahlung, and that the radio emission is produced by nonthermal power-law synchrotron emission (Drake et al. 1992). The correlations imply the co-existence of thermal and nonthermal electrons. In dense interstellar shocks, thermal bremsstrahlung is thought to be the major emission mechanism (Neufeld \\& Hollenbach 1996), but the observed spectral index of -0.7 in radio bands (Curiel et al. 1993), which exceeds the limit of the thermal bremsstrahlung spectral index, implies that nonthermal electrons may also present in the plasma. In solar flares, the radio spectra and x-ray spectra indicate that the sources are hybrid thermal-nonthermal plasmas (Benka \\& Hollman 1992, 1994). In some other astrophysical sources with a high Thomson depth such as x-ray sources (Crider et al. 1997) and gamma ray bursters (Tarvani 1996, Liang 1997), the observed spectra suggest that the sources are hybrid thermal-nonthermal plasmas. Such broad presence of the hybrid plasma makes it important for us to study its emission and absorption processes and the emergent self-absorbed spectrum in the general case. The major emission processes in the hybrid plasma are gyrosynchrotron and bremsstrahlung. The gyrosynchrotron emission includes the cyclotron emission by the nonrelativistic electrons and the synchrotron emission by the relativistic electrons. The bremsstrahlung includes ion-electron bremsstrahlung, pair bremsstrahlung and electron-electron bremsstrahlung. But the electron-electron bremsstrahlung is weaker than the ion-electron bremsstrahlung at low energies (Haug 1985), which is of most concern in this paper for the comparison with gyrosynchrotron processes. Therefore, we will neglect the electron-electron bremsstrahlung. Because of the complexity of calculating the emissivities of the two emission processes, we will confine our discussion to a certain physical regime which will simplify the calculation but still covers a wide range of astronomical phenomena. The magnetic field is assumed to be less than $10^{10} $ Gauss so that no quantum effects have to be considered. We will not consider the Compton scattering, so the Thomson depth $\\tau_T \\ll 1$. We will assume that the thermal electrons are nonrelativistic, or $kT_e \\ll m_ec^2$. The size of the source R is small enough that the photon escaping time $\\frac{R}{c}$ is much less than the time scale of the electron distribution evolution. So the source can be treated as a steady-state one. The formulas we will review for the emissivities are applicable to arbitrary electron distribution functions. The nonthermal fraction $F_{nth} $, which is defined as the fraction of electrons with $\\gamma > \\gamma_{th}$, can be any value between 0 and 1. However, to be a meaningful hybrid source, it must includes enough thermal electrons, otherwise no thermal peak is seen in the particle distribution function. Therefore, we will constrain our discussions to $F_{nth} < 10\\%$. We also assume that the source has a very small fraction of positrons. The combination of gyrosynchrotron and bremsstrahlung emissions by the hybrid plasma can produce some distinct spectra which already have some useful applications (Benka \\& Holman 1992, 1994). The gyrosynchrotron emissions by both thermal electrons and nonthermal electrons explain the harmonic features at the cyclotron frequencies (Benka \\& Holman 1992), which are useful in determining the magnetic field of the sources. And also, the bremsstrahlung spectra by the hybrid electrons fit the soft and hard x-ray spectra of solar flares very well (Benka \\& Holman 1994). However, gyrosynchrotron or bremsstrahlung alone is not enough to explain another spectrum feature of solar flares called spectrum flattening (Lee, Garry \\& Zirin 1994). Ramaty \\& Petrosian (1972) has suggested that power-law synchrotron emissions absorbed by thermal bremsstrahlung can produce the flat spectra. They showed that the solar conditions favor such mechanism. However, the spectrum flattening should not be a unique property of solar flares. The flattened spectra should also be visible in other hybrid sources. The general conditions have to be found that favor such a flat spectrum. To find such conditions is equivalent to determine which, gyrosynchrotron or bremsstrahlung, dominates in the hybrid plasma. In this article, We will calculate the emissivities of gyrosynchrotron and bremsstrahlung emissions to determine the dominant processes and their conditions. We will also discuss the harmonic features and the conditions under which the harmonic features exist. In section 2, we will review the formulas for calculating the emissivities and the absorption coefficients of gyrosynchrotron and bremsstrahlung emissions. In section 3, we will show the general features of the emerging spectra from a hybrid plasma and the dominant emission mechanisms. In section 3, we will find the conditions for the spectrum flattening in a hybrid plasma. In section 4, we will discuss the harmonic features. ", + "conclusions": "" + }, + "9802/astro-ph9802309_arXiv.txt": { + "abstract": "Rotation curves and velocity dispersion profiles are presented for both the stellar and gaseous components along five different position angles (P.A.=$5^{\\circ}$, $50^{\\circ}$, $95^{\\circ}$, $125^{\\circ}$, and $155^{\\circ}$) of the nearby barred spiral NGC~6221. The observed kinematics extends out to about $80''$ from the nucleus. Narrow and broad-band imaging is also presented. The radial profiles of the fluxes ratio \\nii\\ ($\\lambda$ 6583.4 \\AA )/\\ha\\ reveal the presence of a ring-like structure of ionized gas, with a radius of about $9''$ and a deprojected circular velocity of about 280 \\kms . The analysis of the dynamics of the bar indicates this ring is related to the presence of an inner Lindblad resonance (ILR) at 1.3 kpc. NGC~6221 is found to exhibit intermediate properties between those of the early-type barred galaxies: the presence of a gaseous ring at an ILR, the bar edge located between the ILR's and the corotation radius beyond the steep rising portion of the rotation curve, the dust-lane pattern, and those of the late-type galaxies: an almost exponential surface brightness profile, the presence of \\ha\\ regions along all the bar, the spiral-arm pattern. It is consistent with scenarios of bar-induced evolution from later to earlier-type galaxies\\footnote{Tables 3, 4, 5, 6, 7 and 8 are only available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via http://cdsweb.u-strasbg.fr/Abstract.html}. ", + "introduction": "\\object{NGC~6221} is a nearby spiral classified as Sbc(s) by Sandage \\& Tammann (1981) and as SBc(s) by de Vaucouleurs et al. (1991). In The Carnagie Atlas of Galaxies (Panel~189) Sandage and Bedke (1994) describe its morphology as semichaotic with two symmetric heavy dust lanes starting from the nucleus and threading through the middle of the opposite thick arms, which begin at the centre. They also stated that although the galaxy is not strongly barred its dust pattern is similar to that of prototype SBb galaxies, characterized by two straight dust lanes. An overview of the optical properties of the galaxy is given in Table~1. NGC~6221 forms an apparent physical pair with the late-type spiral \\object{NGC~6215}. They have a systemic velocity of $1465\\pm10$ \\kms\\ (this paper) and $1521\\pm43$ \\kms\\ (RC3) respectively. Their angular distance is about 26$'$ (RC3) corresponding to a projected linear separation of about 227 kpc at a mean distance of 30 Mpc ($H_0 = 50\\;\\rm km\\;s^{-1}\\;Mpc^{-1}$) The galaxy is also possibly interacting with two newly discovered low-surface brightness galaxies nearby (Koribalski 1996a). The dynamics of the ionized gas component in this galaxy has been studied by Pence \\& Blackman in 1984 (hereafter P\\&B). They measured a conspicuous S-shaped pattern in the ionized gas velocity field, showing that large velocity gradients occur at the position of the dust lanes in the bar. They have been interpreted as the signature of shock fronts of the gas which reverses its motion from outward to inward as it passes through them. Philipps (1979) reported that the nuclear spectrum of NGC~6221 exhibits signatures of emission originating from \\hii\\ regions and also a weak Seyfert~2 component. He concluded however that the emission line spectrum can be interpreted as gas being ionized by hot young stars without the need to introduce a non-thermal component in the galaxy nucleus. The analysis of the spatial distribution of the ionized gas (Durret \\& Bergeron 1987) confirmed this picture. Some dozens of \\hii\\ regions were identified in the disk region extending as far as 9~kpc from the galaxy nucleus. Although the emission in the nucleus appears to be extended ($\\sim 1 \\; {\\rm kpc}^2$), the authors conclude that the main ionizing sources are OB stars and that UV and X-ray radiation from the nuclear source might not penetrate the heavily obscured nuclear region. Dottori et al. (1996) also recently suggested that the nucleus of NGC~6221 may harbour sources of type \\hii\\ and also a Seyfert~2 source. In this paper we study the kinematics of the ionized gas and the stellar components in NGC~6221. We present the velocity curves and the velocity dispersion profiles of gas and stars obtained along five different position angles by means of long-slit spectroscopy. \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{ds7406f01.eps}} \\caption{\\ha$+$\\nii\\ emission lines image of NGC~6221. The slit positions (P.A.=$5^{\\circ}$, $50^{\\circ}$, $95^{\\circ}$, $125^{\\circ}$, and $155^{\\circ}$) of the obtained spectra are plotted. The {\\it black bars\\/} show the position and the size of the regions (along the different slit positions) where the \\nii/\\ha\\ fluxes ratio is greater than 0.5. The image orientation and scale are indicated } \\label{fig:halpha} \\end{figure} ", + "conclusions": "The possible link between bar component and nuclear activity has a long history either in the observational and in the theoretical field (e.g. Shlosman 1994). Recent observations have lead Ho et al. (1997) to disagree with this idea and to conclude that, even though the presence of a bar can enhance star formation in the bulge region, it does not affect significantly the active nucleus. Our data for NGC~6221 support this interpretation, because the physical properties of the central line emitting region indicate clearly that ionization comes from the presence of stellar sources. An AGN and its observational signatures might be obscured by gas and dust, but this is, of course, speculation. The asymmetries of the velocity curves outside the bar region may be due to a tidal distorsion caused by NGC~6215. Koribalski (1996a) discovered an \\hi\\ bridge between NGC~6221 and NGC~6215 and later (Koribalski 1996b) suggested a possible interaction of NGC~6221 with the two newly discovered low-surface brightness galaxies nearby, namely \\object{BK1} and \\object{BK2}. Elmegreen \\& Elmegreen (1989) and later Keel (1996) proposed a dichotomy of bar properties between early-type (Sbc and earlier) and late-type galaxies (Sc and later) based on observational properties such as bar size, density profile of the bar, star formation, and so on. Early-type bars tend to be a dominant component in the galaxy with a flat density profile extending out to the corotation radius, where they drive symmetric spiral arms (Combes \\& Elmegreen 1993). On the contrary, bars in late-type galaxies are small with exponentially decreasing density profiles, which extend only to the ILR radius. These kind of bars are unable to influence the wave pattern of the stellar disk. This view is also supported by the model of Combes \\& Elmegreen (1993). NGC~6221 is found to exhibit intermediate properties between early-type and late-type barred spiral galaxies, as discussed below: \\begin{description} \\item [{\\it i) Late-type features:\\/}] The surface brightness profile shows an exponential decrease; the star formation, i.e. the \\ha\\ emission, is detected along all the bar; the arm structure is typical of a late-type. \\item [{\\it ii) Early-type features:\\/}] The edge of bar is located between the outer ILR radius and the corotation radius beyond the rising part of the velocity curve; the strong gradient of gas velocity curve in the central $2$ kpc; the dust-lane pattern is similar to that of the prototype SBb galaxies (see \\object{NGC~1300} in Panels S8 and 154 in the CAG as example); the radial profiles of the fluxes ratio of the \\ha\\ and \\nii\\ emission lines reveal the presence of a ring structure associated with an ILR; moreover NGC~6221 is interacting with NGC~6215, and Elmegreen et al. (1990) have demonstrated the presence of bars in paired early-type spirals. \\end{description} Indeed the case of NGC~6221 is analogous to that of certain other galaxies, notably NGC~1300 (Combes \\& Elmegreen 1993). It has been classified as a late-type barred spiral because of the presence of gas, star formation and spiral arms, but it could be considered as an early-type due to the dynamical properties of its bar. Pfenniger (1992) has suggested that a major effect due to the presence of a bar could be that the galaxy evolves towards earlier morphological types. Our study indicates that NGC~6221 is a typical case of a late-type barred spiral evolving to an earlier type. If our interpretation is correct, it is understandable that there is so difficult to reach agreement about its morphological classification which is Sbc(s) in RSA and SBc(s) in RC3. Sometimes the criteria of classification in the Hubble sequence cannot describe exactly the complicated scenario of the evolution of galaxies, as appears to be the case for NGC~6221." + }, + "9802/astro-ph9802079_arXiv.txt": { + "abstract": "The recent discovery of two flux decrements in deep radio maps obtained by the VLA and the Ryle Telescope can have powerful implications for the density parameter of the Universe, $\\Omo$. We outline these implications by modeling the decrements as the thermal Sunyaev--Zel'dovich (SZ) effect from two clusters {\\em assuming} their properties are similar to those of the low redshift population. In this case, the absence of any optical or X--ray counterparts argues that the clusters must be at large redshifts. We highlight the difficulty this poses for a critical cosmology by a comparison with a fiducial open model with $\\Omo=0.2$ ($\\lamo=0$). Applying the phenomenological X--ray luminosity--temperature relation needed to explain the EMSS cluster redshift distribution, as inferred by Oukbir and Blanchard (1997), we convert the X--ray band upper limits to {\\em lower} limits on the clusters' redshifts. Comparison of the counts implied by these two SZ detections with model predictions, for clusters with redshifts larger than these lower limits, illustrates quantitatively the inability of the critical cosmology to account for such high--redshift clusters. On the other hand, the open model with $\\Omo=0.2$ remains consistent with the existence of the two objects; it possibly has, however, difficulties with current limits on spectral distortions and temperature fluctuations of the cosmic microwave background. The discussion demonstrates the value of SZ cluster searches for testing cosmological models and theories of structure formation. ", + "introduction": "Most favored theories of structure formation in the Universe are based on the gravitational growth of initially small density perturbations with Gaussian statistics. The Gaussian characteristic finds its way into the mass function of cosmic structures and leads to the expectation of an exponentially rapid decline in the number density of objects with increasing mass (Press \\& Schechter \\cite{PS}). The interesting consequence is that the abundance of massive objects, such as galaxy clusters, is then extremely sensitive to the power spectrum (amplitude and shape) of the density fluctuations and their growth rate with redshift. Since the growth rate is controlled by the density parameter of the Universe, $\\Omo$, this means that the {\\em shape} of the redshift distribution of clusters of a given mass is sensitive to $\\Omo$. In fact, once constrained by local data, the redshift distribution depends {\\em only} on the underlying cosmology, i.e., $\\Omo$ (and to a lesser extent, on the cosmological constant $\\Lambda$). In other words, the redshift distribution provides a probe of the density of the Universe (Oukbir \\& Blanchard \\cite{OB0}, \\cite{OB}; Blanchard \\& Bartlett \\cite{b2}). The sense is such that we expect many more clusters at large redshift if $\\Omo<1$, because the growth rate is suppressed by the rapid expansion at late times in open models, leading to less evolution towards the past. Essential for the application of this probe of $\\Omo$ is the existence of an easily observed cluster quantity well correlated with virial mass. It has been cautioned by many authors that the velocity dispersion of cluster member galaxies is too easily inflated by contamination of interlopers along the line--of--sight (e.g., Lucy \\cite{interlop1}; Frenk et al. \\cite{interlop2}; Bower et al. \\cite{interlop3}). Many have turned instead to the X--ray temperature of the intracluster medium (ICM). On theoretical grounds, it is believed that the gas is heated by infall to the virial temperature of the cluster gravitational potential well. Numerical simulations in fact support the existence of a tight relation between virial mass and X-ray, which is to say, emission weighted, temperature (Evrard \\cite{Tsim1}; Evrard et al. \\cite{Tsim2}). The dependence of the relation is as expected based on the idea that the gas is shock heated to the virial temperature on infall, although the simulations indicate that there is an incomplete thermalization of the gas, resulting in a temperature slightly ($\\sim 20\\;$ \\%) smaller than the virial value. The X--ray temperature is to be preferred over the X--ray luminosity as an indicator of virial mass, because the X--ray luminosity depends not only on the temperature, but also on the quantity of gas and on its density, or, what is equivalent, the spatial distribution of the ICM. This spatial distribution is difficult to model, particularly because there is at present no understanding of the origin of the ICM core radius. Use of the X--ray temperature function to constrain models of structure formation is rather well developed as a subject (e.g., see Bartlett \\cite{casa} and references therein). Temperature data on clusters at $z>0$ is just now becoming available, and the possibility of even higher $z$ data from future space missions like XMM makes the application of the redshift distribution test proposed by Oukbir \\& Blanchard a real possibility over the near term (see, e.g., Sadat et al. \\cite{SBO}). The Sunyaev--Zel'dovich (SZ) (Sunyaev \\& Zel'dovich \\cite{SZ}) effect offers another, complimentary approach to the problem of applying mass function evolution as a probe of $\\Omo$. Due to the distance independence of the surface brightness of the distortion, the effect represents an efficient method of finding high redshift clusters. This should be contrasted with X--ray emission, whose surface brightness suffers the $(1+z)^{-4}$ cosmological dimming. Moreover, as will be developed below, the SZ effect has other, important advantages over X--ray studies: the integrated SZ signal of a cluster, its flux density (measured in Jy), is proportional to the {\\em total hot gas mass} times the {\\em particle weighted} temperature. This means that the signal is independent of the gas' spatial distribution and that the temperature involved is closely tied to the cluster virial mass, by energy conservation during collapse, for it is simply the total energy of the system divided by the number of gas particles. For the same reason, this temperature should also be {\\em much less sensitive to any temperature structure in the gas} than the X--ray (emission weighted) temperature. Thus, the SZ effect is an observable which combines ease of theoretical modeling with ease of detection at large $z$. All of this has prompted several calculations of the expected SZ number counts and redshift distribution of SZ selected clusters, and their dependence on the cosmological parameters and ICM evolution (Korolyov et al. \\cite{SZcounts1}; Bartlett \\& Silk \\cite{SZcounts2}; Markevitch et al. \\cite{SZcounts3}; Barbosa et al. \\cite{SZcounts4}; Eke et al. \\cite{Tcon8}; Colafrancesco et al. \\cite{SZcounts5}). The future of this kind of study appears bright with the prospect of the Planck Surveyor satellite mission ({\\tt http://astro.estec.esa.nl/Planck/}). Ground based efforts have also made astounding progress recently, and it is already feasible to map $\\sim 1$ square degree of sky to produce number counts down to flux levels sufficient to test theories (Holzapfel, private communication). In this paper, we discuss what {\\em may} already be an indication of clusters at very large redshift and the resulting implications. We refer to two radio decrements, one found in a deep VLA field (Richards et al. \\cite{vla:radio}) and the other detected by the Ryle Telescope during an observation of a double quasar system (Jones et al. \\cite{ryle:radio}). Although these detections await definitive confirmation, we will nevertheless proceed to outline here the implications of their explanation as the thermal SZ effect produced by two clusters. What makes just two such objects of great importance is the fact that no optical or X--ray counterparts have been observed, and the flux limits in the X--ray are so stringent that the clusters would have to be at large redshift (Richards et al. \\cite{vla:radio}; Kneissl \\cite{kneissl1}; Kneissl et al. \\cite{kneissl2}). This is of paramount importance because, as we have mentioned, massive clusters (say $M \\geq 10^{15} $M$_{\\odot}$) at large $z$ are not expected in critical models. Our goal in this paper is to quantify just how badly critical models fare in this regard. We emphasize that the modeling is based on the observed characteristics of the galaxy cluster population, in particular the X--ray luminosity--temperature relation and constraints on its potential evolution. This excludes from the present discussion the possibility of a large class of low luminosity clusters (both optical and X--ray). We believe that a clear discussion in this restricted context is nevertheless useful. (For this reason we dubb this work a ``tale''!). The procedure also demonstrates the great potential of SZ cluster searches for constraining theories of structure formation. The plot of the tale proceeds as follows: In the next section, we introduce the two radio decrements and their properties which will be used later. Then we outline our modeling of the SZ cluster population and of the two radio decrements. This is followed by a discussion of the X--ray emission to be expected from clusters producing the observed SZ signals and the {\\em minimum} redshifts imposed by the X--ray flux upper limits; this represents a key element of our tale. Finally, we discuss the results and various caveats in the analysis before bringing an end to the tale with a brief summary. ", + "conclusions": "The principal result of this work is given in Figure \\ref{SZcounts}, where we see clearly and quantitatively the difficulty faced by a critical cosmology if the VLA and RT radio decrements are representative of the cluster population's SZ effect. We also show two other observational constraints in this figure. The SuZIE instrument recently reported no detections down to 12 mJy (at $\\lambda=0.75$ mm) on blank sky covering $\\Omega^{\\rm SuZIE}_{tot}=0.06\\;$ deg$^2$. The resulting upper limit is shown in the figure as the downward--pointing triangle between the VLA and RT boxes. The rightmost box presents one possible interpretation of the results of the OVRO RING experiment (Myers et al. \\cite{RING}). In this experiment, there is one field representing a $>5\\;\\sigma$ fluctuation and for which no source has yet been identified. If we suppose that this fluctuation is due to the SZ effect, then we deduce the constraints given by the box based on this detection, at $\\sim 60$ mJy ($\\lambda=0.75$ mm), over a total survey area of $\\Omega_{\\rm obs}^{\\rm RING}=0.1\\;{\\rm deg^2}$. Once again, we give the 95\\% {\\em one--sided} confidence limits, and the range in flux density is from the detected flux density to the corrected, total flux density for $z=1$. Alternatively, we could use the RING data as an upper limit to the source counts, arguing that this one fluctuation is not the result of a cluster. We then obtain the upper limit given as the downwards pointing arrow at the far left-hand side of the RING box. In contrast to the critical model, our fiducial open model has no difficulty in accounting for all of the constraints shown in the figure. It is the strong upper limits on any X--ray flux from the objects, forcing them to be at large redshift (Figure 3), that makes the distinction between the two models not just a question of a factor of a few, but of {\\em orders of magnitude}. The exponential behavior of the mass function provides us with large ``leverage'' to discriminate models at high $z$. To facilitate the presentation in the figure, we have used a lower limit of $z>1$ on the critical model, but it should be emphasized that the actual limit from Figures 3 and 4 is at least $z>2$ (for the RT), and this would put the curve completely off to the bottom left of the plot. Thus, a straightforward interpretation of the results is that the critical model is ruled out. Let us now discuss the various caveats to this `straightforward' interpretation. The first thing that must again be emphasized is that the line of argument relies heavily on the idea that the radio decrements (are real) are due to gas heated to the virial temperature of collapsed objects and that these objects behave in a manner similar to what is known about X--ray clusters. It is perhaps possible that the decrements are not due to such objects. In this case, a different approach than the one presented here is needed. The calculation in this paper was made in order to understand what the radio decrements imply if one wishes to explain them by what we know as galaxy clusters. Even within this context, there are several issues we should address. The first is the scatter in the $L(T,z)$ relation. Arnaud and Evrard (1998) have shown that this scatter is intrinsically smaller than previously thought, by considering a cluster sample with high quality X--ray observations. The scatter is seen to be $~20\\%$ around a relation $\\propto T^{2.88}$ for non--cooling flow clusters. There are, however, a couple of clusters below the mean relation and well outside the scatter. It is possible that the SZ detections are the first indication of a larger than expected population of intrinsically underluminous clusters. Next, as already remarked, we chose for simplicity to hold the ICM gas mass fraction, $\\fgas$, constant over mass and redshift; one could imagine that it is in fact a function of both. Colafrancesco et al. (\\cite{SZcounts5}) have investigated the SZ source counts including the possible effects of cluster evolution. Accounting for cluster evolution in the present context, however, will not substantially change the conclusion in respect to the critical model -- we are already using, in this model, a large gas mass fraction, as supported by X--ray observations (Evrard \\cite{gasfrac}) (This fraction is in violation of primordial nucleosynthesis predictions for $\\Omo=1$, another problem for the critical model [White et al. \\cite{baryoncris}]). Any reasonable evolution would thus cause this fraction to decrease with either mass or redshift, or both, thereby {\\em decreasing} the counts predicted by the critical model -- things would get worse. One thing which could help the critical model is if the intergalactic medium surrounding the virialized region of a cluster was heated to close to 1 keV. One possible mechanism could be the diffusion of electrons through the shock front that heats the gas to the virial temperature (Chi\\`eze et al. 1997). This would increase the SZ flux density associated with a given cluster mass, pushing the predicted curves to the right. Such a mechanism could be in operation around clusters, but the factor needed to reconcile the critical model with the indicated counts seems unreasonable. A particular source of concern is the fact that the RT object was found by observing known quasar systems, and not a priori blank fields. This seems to be less of a worry for the VLA fields, because no double quasar systems were known prior to the observations (although one was subsequently found...). We can get a feeling, at least, for the effect of a possible bias on our results in the following manner: Supposing that instead of representing random fields on the sky for a cluster search, double quasar systems are {\\em always} associated with clusters (which are responsible, say, for the two images). Then another way to proceed with the RT detection would be to take the observed sky density of such double quasar systems as the SZ source counts. Over a survey area of $\\sim 60\\; {\\rm deg}^2$, Schneider et al (\\cite{quasdens}) found 90 quasars, 3 of which are close double systems (with separations less than $\\sim 400\\;$~arcsecs). One of these three is in fact PC1643+4631, i.e., the RT field with the radio decrement. The implied counts are then $\\sim 0.05\\; {\\rm deg}^{-2}$, with {\\em one--sided 95\\% confidence} upper and lower limits of $\\sim 0.13\\; {\\rm deg}^{-2}$ and $\\sim 0.013\\; {\\rm deg}^{-2}$, respectively. Looking at Figure 5, we see that this will not drastically alter the severity of the critical model's difficulty (remember that $z>2$ really applies to the RT object)." + }, + "9802/astro-ph9802286_arXiv.txt": { + "abstract": "The old question of rotational braking of Ap Si stars is revisited on the empirical side, taking advantage of the recent Hipparcos results. Field stars with various evolutionary states are considered, and it is shown that the loose correlation between their rotational period and their surface gravity is entirely compatible with conservation of angular momentum. No evidence is found for any loss of angular momentum on the Main Sequence, which confirms earlier results based on less reliable estimates of surface gravity. The importance of reliable, fundamental $T_{\\rm eff}$ determinations of Bp and Ap stars is emphasized. ", + "introduction": "It is well known that chemically peculiar stars of the Ap and Am types are rotating slower than their normal counterparts (e.g. North 1994). The question then arises, whether slow rotation is acquired during the main sequence life of the star, or before its arrival on the ZAMS, i.e. during the proto-stellar phase. Havnes \\& Conti (1971) had suggested that magnetic stars undergo magnetic braking during their main sequence lifetime, due to mass accretion from the interstellar medium, while Strittmatter \\& Norris (1971) proposed the same, but due to mass loss. These theoretical considerations seemed to get support from observational evidence when Wolff (1975, 1981), Stift (1976) and Abt (1979) found some correlation between the radii or ages of Ap stars and their rotational period obtained from their photometric or spectroscopic variation. On the contrary, Hartoog (1977) concluded that magnetic Ap stars in young clusters do not rotate faster than those in older clusters, and this conclusion was also reached by North (1984a, b, 1985, 1986, 1987), Borra et al. (1985) and Klochkova \\& Kopylov (1985). The apparent correlation between radius and rotational period has been commented by Hensberge et al. (1991), who conclude that this correlation is real but possibly due to a detection bias depending on the inclination angle, and by Stepien (1994), who concluded on the contrary that this correlation does simply not exist, if spurious rotational periods are duly excluded. Using the $\\log P_{\\mathrm rot}$ vs. $\\log g$ diagram for field stars, North (1985, 1986, 1992) showed that for Si stars, there is indeed a trend towards longer periods for low-gravity stars, but which can be entirely explained by conservation of angular momentum as the star evolves with increasing radius within the main sequence. In this note, we revisit the $\\log P_{\\mathrm rot}$ vs. $\\log g$ diagram for field stars having both a rotational period in the literature and a reliable surface gravity, the latter being either spectroscopic or obtained from Hipparcos data. ", + "conclusions": "New surface gravities of magnetic Bp and Ap stars obtained from the Hipparcos parallaxes, as well as homogeneous spectroscopic gravities, have been used to reconsider how the rotational period of such stars varies with age. The result is entirely consistent with previous works suggesting that field Si stars do not undergo any significant magnetic braking during their life on the Main Sequence; it is also more firmly based than earlier studies made on field Ap stars. Therefore, the slow rotation of these objects must be a property acquired {\\it before} they arrive on the ZAMS. How this occurs has just been explored by Stepien (1998) but further investigations remain worthwhile. On the other hand, this study has shown that $\\log g$ values obtained from Hipparcos luminosities may be overestimated by up to 0.2 dex for some extreme Ap stars, probably through an overestimate of their $T_{\\rm eff}$. This shows how badly fundamental determinations of $T_{\\rm eff}$ are needed for these stars." + }, + "9802/astro-ph9802214_arXiv.txt": { + "abstract": "We study the effect of tidal torques on the collapse of density peaks through the equations of motion of a shell of barionic matter falling into the central regions of a cluster of galaxies. We calculate the time of collapse of the perturbation taking into account the gravitational interaction of the quadrupole moment of the system with the tidal field of the matter of the neighbouring proto-clusters. We show that within high-density environments, such as rich clusters of galaxies, tidal torques slow down the collapse of low-$\\nu $ peaks producing an observable variation in the time of collapse of the shell and, as a consequence, a reduction in the mass bound to the collapsed perturbation. Moreover, the delay of the collapse produces a tendency for less dense regions to accrete less mass, with respect to a classical spherical model, inducing a {\\it biasing} of over-dense regions toward higher mass. Finally we calculate the bias coefficient using a selection function properly defined showing that for a Standard Cold Dark Matter (SCDM) model this ${\\it bias}$ can account for a substantial part of the total bias required by observations on cluster scales. ", + "introduction": "It has long been speculated on the fundamental role that the angular momentum could play in determining the fate of collapsing proto-structures and several models have been proposed to correlate the galaxy type with the angular momentum per unit mass of the structure itself (Faber 1982; Kashlinsky 1982; Fall 1983). Some authors (see Barrow \\& Silk 1981; Szalay \\& Silk 1983 and Peebles 1990) have proposed that non-radial motions would be expected within a developing proto-cluster due to the tidal interaction of the irregular mass distribution around them, typical of hierarchical clustering models, with the neighbouring proto-clusters. The kinetic energy of these non-radial motions prevents the collapse of the proto-cluster, enabling the same to reach statistical equilibrium before the final collapse (the so-called previrialization conjecture by Davis \\& Peebles 1977, Peebles 1990). This effect may prevent the increase in the slope of the mass autocorrelation function at separations given by $\\xi(r,t) \\simeq 1$, expected in the scaling solution for the rise in $\\xi(r,t)$ but not observed in the galaxy two-point correlation function. The role of non-radial motions has been pointed by several authors (see Davis \\& Peebles 1983: Gorski 1988; Groth et al. 1989; Mo et al. 1993; van de Weygaert \\& Babul 1994; Marzke et al. 1995 and Antonuccio-Delogu \\& Colafrancesco 1995). Antonuccio-Delogu \\& Colafrancesco derived the conditional probability distribution $f_{pk}({\\bf v}| \\nu)$ of the peculiar velocity around a peak of a Gaussian density field and used the moments of the velocity distribution to study the velocity dispersion around the peak. They showed that regions of the proto-clusters at radii, $ r$, greater than the filtering length, $ R_{f}$, contain predominantly non-radial motions. \\\\ Non-radial motions change the energetics of the collapse model by introducing another potential energy term. In other words one expects that non-radial motions change the characteristics of the collapse and in particular the {\\it turn around} epoch, $t_{m}$, and consequently the critical threshold, $ \\delta_{c}$, for collapse. Here, we want to remind that $ t_{m}$ is the time at which the linear density fluctuations, that generate the cosmic structures, detach from the Hubble flow. The turn-around epoch is given by: \\begin{equation} t_{m} = \\left[\\frac{ 3 \\pi}{32 G \\rho_{b}} ( 1 +\\overline{\\delta}) \\right]^{1/2} (1+z)^{3/2} \\end{equation} where $ \\rho_{b} $ is the mean background density, $z$ is the redshift and $\\overline{\\delta}$ is the mean over-density within the non-linear region. After the {\\it turn around} epoch, the fluctuations start to recollapse. As known for a spherical top hat model, the perturbation of the density field is completely collapsed when \\begin{equation} \\overline{\\delta} = \\delta_{c}=(3/5)(\\frac{3 \\pi T_{c}}{4 t_{m}})^{2/3} = 1.68 \\end{equation} where $T_{c}$ is the time of collapse which is twice the turn around epoch. One expects that non-radial motions produce firstly a change in the turn around epoch, secondly a new functional form for $ \\delta_{c}$, thirdly a change in the mass function calculable with the Press-Schechter (1974) formula and finally a modification of the two-point correlation function. As we shall show in a forthcoming paper (Del Popolo \\& Gambera 1997b) non-radial motions can reduce several discrepancies between the SCDM model and observations: the strong clustering of rich clusters of galaxies ($\\xi _{cc}(r) \\simeq (r/25h^{-1}Mpc) ^{-2}$) far in excess of CDM predictions (Bahcall \\& Soneira 1983), the X-ray temperature distribution function of clusters over-producing the observed cluster abundances (Bartlett \\& Silk 1993).\\\\ For the sake of completeness, we remember that alternative models with more large-scale power than SCDM have been introduced in order to solve the latter problem. Several authors (Peebles 1984; Efstathiou et al. 1990; Turner 1991; White et al. 1993) have lowered the matter density under the critical value ($% \\Omega _m<1$) and have added a cosmological constant in order to retain a flat universe ($\\Omega _m+\\Omega _\\Lambda =1$) . The spectrum of the matter density is specified by the transfer function, but its shape is affected because of the fact that the epoch of matter-radiation equality is earlier, $1+z_{eq}$ being increased by a factor $1/\\Omega_{m}$. Around the epoch $z_\\Lambda $ the growth of the density contrast slows down and ceases after $z_\\Lambda $. As a consequence the normalisation of the transfer function begins to fall, even if its shape is retained.\\\\ Mixed dark matter models (MDM) (Bond et al. 1980; Shafi \\& Stecker 1984; Valdarnini \\& Bonometto 1985; Holtzman 1989; Schaefer 1991; Schaefer \\& Shafi 1993; Holtzman \\& Primack 1993) increase the large-scale power because free-streaming neutrinos damp the power on small scales. Alternatively changing the primeval spectrum several problems of SCDM are solved (Cen et al. 1992). Finally, it is possible to assume that the threshold for galaxy formation is not spatially invariant but weakly modulated ($2\\%-3\\%$ on scales $% r>10h^{-1}Mpc$) by large scale density fluctuations, with the result that the clustering on large-scale is significantly increased (Bower et al. 1993). \\\\ Moreover, this study of the role of non-radial motions in the collapse of density perturbations can help us to give a deeper insight in to the so-called problem of biasing. As pointed out by Davis et al. (1985), unbiased CDM presents several problems: pairwise velocity dispersion larger than the observed one, galaxy correlation function steeper than that observed (see Liddle \\& Lyth 1993 and Strauss \\& Willick 1995). The remedy to these problems is the concept of biasing (Kaiser 1984), i.e. that galaxies are more strongly clustered than the mass distribution from which they originated. The physical origin of such biasing is not yet clear even if several mechanisms have been proposed (Rees 1985; Dekel \\& Silk 1986; Dekel \\& Rees 1987; Carlberg 1991; Cen \\& Ostriker 1992; Bower et al. 1993; Silk \\& Wyse 1993). Recently Colafrancesco, Antonuccio-Delogu \\& Del Popolo (1995, hereafter CAD) and Del Popolo \\& Gambera (1997a) have shown that dynamical friction delays the collapse of low-$\\nu$ peaks inducing a bias of dynamical nature. Because of the dynamical friction, under-dense regions in clusters (the clusters outskirts) accrete less mass than that accreted in absence of this dissipative effect and as a consequence over-dense regions are biased toward higher mass (Antonuccio-Delogu \\& Colafrancesco 1994 and Del Popolo \\& Gambera, 1996). Non-radial motions acts in a similar way to dynamical friction: they delay the shell collapse consequently inducing a dynamical bias similar to that produced by dynamical friction. This dynamical bias can be evaluated defining a selection function similar to that given in CAD and using Bardeen, Bond, Szalay \\& Kaiser (1986, hereafter BBKS) prescriptions. \\\\ The methods used in this paper are fundamentally some results of the statistics of Gaussian random fields, the biased galaxy formation theory and the spherical model for the collapse of density perturbations. In particular, we calculate the specific angular momentum acquired by protoclusters and the time of collapse of protoclusters using the Gaussian random fields theory and the spherical collapse model following Ryden's (1988a, hereafter R88a) approach. The selection function that we introduce is general and obtained by the only hypothesis of Gaussian density field. The approach and the final result is totally different from BBKS selection function and similar to that of Colafrancesco, Antonuccio \\& Del Popolo (1995). Only the biasing parameter is obtained from a BBKS approximated formula. This choice will be clarified in the following sections of the paper.\\\\ The plan of the paper is the following: in Sect. 2 we obtain the total specific angular momentum acquired during expansion by a proto-cluster. In Sect. 3 we use the calculated specific angular momentum to obtain the time of collapse of shells of matter around peaks of density having $\\nu_c = 2, 3, 4$ and we compare the results with Gunn \\& Gott's (1972, hereafter GG) spherical collapse model. In Sect. 4 we derive a selection function for the peaks giving rise to proto-structures while in Sect. 5 we calculate some values for the bias parameter, using the selection function derived, on three relevant filtering scales. Finally in Sect. 6 we discuss the results obtained. ", + "conclusions": "In this paper we have studied the role of non-radial motions on the collapse of density peaks solving numerically the equations of motion of a shell of barionic matter falling into the central regions of a cluster of galaxies. We have shown that non-radial motions produce a delay in the collapse of density peaks having a low value of $ \\nu$ while the collapse of density peaks having $ \\nu > 3$ is not influenced. A first consequence of this effect is a reduction of the mass bound to collapsed perturbations and an increase in the critical threshold, $ \\delta_{c}$, which now is larger than that of the top-hat spherical model and depends on $ \\nu$. This means that shells of matter of low density have to be subject to a larger gravitational potential, with respect to the homogeneous GG's model, in order to collapse. \\\\ The delay in the proto-structures collapse gives rise to a dynamical bias similar to that described in CAD whose bias parameter may be obtained once a proper selection function is defined. The selection function found is not a pure Heaviside function and is different from that used by BBKS to study the statistical properties of clusters of galaxies. Its shape depends on the effect of non-radial motions through its dependence on $\\delta_{c}(\\nu)$. The function $t(\\nu)$ selects higher and higher density peaks with increasing value of $ R_{f}$ due to the smoothing effect of the density field produced by the filtering procedure. Using this selection function and BBKS prescriptions we have calculated the coefficient of bias $b$. \\\\ On clusters scales for $R_{f} = 4h^{-1}Mpc$ we found a value of $b= 2.25$ comparable with that obtained from the mean mass-to-light ratio of clusters, APM survey, or from N-body simulations combined with hydrodynamical models (Frenk et al. 1990). Morever, the value of the coefficient of biasing $ b$ that we have calculated is comparable with the values of $ b$ given by Kauffmann et al. (1996). This means that non-radial motions play a significant role in determining the bias level. In our next paper (Del Popolo \\& Gambera, in preparation) we make a more detailed analysis on the problem of the bias. \\begin{flushleft} {\\it Acknowledgements} \\end{flushleft} We would like to thank the anonymous referee for insightful comments which led us to explain better our ideas. Besides, we are grateful to V. Antonuccio-Delogu for helpful and stimulating discussions during the period in which this work was performed." + }, + "9802/astro-ph9802022_arXiv.txt": { + "abstract": "We report on new RXTE observations to search for millisecond X-ray variability from the bright Galactic bulge X-ray source GX 3+1. Although kilohertz quasiperiodic oscillations (QPO) have been detected now in 14 low mass X-ray binary (LMXB) systems they have, somewhat mysteriously, not yet been seen in the bright, persistent, Galactic bulge sources; GX 3+1, GX 9+1, GX 9+9 and GX 13+1. These systems have been typically classified as atoll sources and when bright are seen on the banana or upper banana branch in the X-ray color - color diagram. During our observations the source was also in a banana state and we did not detect kHz QPO. We place an upper limit on the amplitude of kHz QPO of about 1\\% rms. ", + "introduction": " ", + "conclusions": "" + }, + "9802/astro-ph9802352_arXiv.txt": { + "abstract": "We analyze bright-end ($K = 10 -17$) galaxy counts from a number of near-infrared galaxy surveys. All studies available as of mid-1997, considered individually or collectively, show that the observed near-infrared galaxy number counts at low magnitudes are inconsistent with a simple no-evolution model. We examine evolutionary effects and a local underdensity model as possible causes of this effect. We find that the data are fit by either a factor of $1.7$ - $2.4$ deficiency of galaxies out to a redshift of $z = 0.10$ - $0.23$, depending on the k corrections and evolution (e-) corrections used and the adopted values of the Schechter luminosity function parameters, or by unexpectedly strong low redshift evolution in the K-band, leading to (e+k)-corrections at $z= 0.5$ that are as much as 60\\% larger than accepted values. The former possibility would imply that the local expansion rate on scales of several hundred Mpc exceeds the global value of $H_0$ by up to 30\\% and that the amplitude of very large scale density fluctuations is far larger than expected in any current cosmogonic scenario. The latter possibility would mean that even the apparently most secure aspects of our understanding of galaxy evolution and spectral energy distributions are seriously flawed. ", + "introduction": "The past ten years have seen a proliferation of near-infrared galaxy surveys providing us with increasingly accurate K-band galaxy counts. The surveys have covered both the bright end (Gardner et al. 1996, Gard96 hereafter; Huang et al. 1997, Huang97 hereafter; Glazebrook et al. 1994; Gardner, Cowie, \\& Wainscoat 1993; Mobasher, Ellis \\& Sharples 1986) and faint magnitudes (Moustakas et al. 1997, Djorgovski et al. 1995; McLoed et al. 1995; Soifer et al. 1994; Jenkins \\& Reid 1991.) Galaxy counts are an essential tool in the study of galaxy evolution and cosmological geometry. The advantage of near-infrared galaxy counts over optical counts is that the former are much less sensitive to stellar population evolution and the K-band k-corrections are smoother and better understood than in the bluer region of the spectrum. In particular, they are much less affected by bursts of star formation and internal extinction by dust, K-band galaxy counts are therefore the ideal tool for probing large scale density variations. Huang97 report that the slope of their bright-end galaxy counts is steeper than that predicted by a no-evolution model. The authors dismiss known observational error and evolution as likely causes for this effect and describe a heuristic model with a local deficiency of galaxies by a factor of 2 on scales sizes of around 300 $h^{-1} Mpc$ as a possible cause. But Gard96 find that this model does not fit their galaxy counts. In this paper, we provide a more in depth study of bright-end galaxy counts in order to determine if other published galaxy counts point to a similar local deficiency of galaxies and to obtain a more precise picture of this underdensity using all K band galaxy counts available as of mid-1997. The possible existence of a very large scale local underdensity merits close scrutiny since it could have profound implications for the determination of $H_{0}$ (Turner, Cen, \\& Ostriker, 1992.) We also attempt to quantify the luminosity evolution that would be needed to account for the count slope discrepancy, the most plausible alternative explanation. The data used, models, and best fit methods are described in sections 2, 3, and 4 respectively. Results are presented in section 5 and discussed in section 6. ", + "conclusions": "Turner et al. (1992) found an approximate expression of the correction from the local to global $H_{0}$, $(\\Delta H_{0}/H_{0}) = -0.6\\times \\Delta n_{gal} \\Omega^{0.4}$, where $\\Delta n_{gal}$ is the over or under density within the local volume. This would lead, in our case to a local evaluation of $H_{0}$ that is as much as $30\\%$ (for $\\Omega \\leq 1$) higher than the global value. However the authors warn that the presence of coherent structure with sizes $> 10000$ km/s might lead to a more extreme effect. Recently, Kim et al. (1997) have shown that $(\\Delta H_{0}/H_{0}) < 0.10$ if $\\Omega_{M} \\leq 1$, using seven supernovae with $0.35 < z < 0.65$. However the authors mention that errors in absolute magnitude calibrations could affect this ratio quite strongly, pointing to the 0.09 mag difference between the absolute magnitude calibrations used in their paper and by Riess, Press, \\& Kirshner (1996) which could lead to a ratio of $(\\Delta H_{0}/H_{0}) < 0.21$. The possibility that the local and the global values of $H_{0}$ differ by of order 20\\% cannot yet be ruled out directly. Zehavi et al. (1998) have analyzed the peculiar velocities of 44 Type Ia supernovae and found a deviation from the Hubble law consistent with a void of $\\sim 20\\%$ underdensity surrounded by a dense wall at $70h^{-1} Mpc.$ This result is consistent with those obtained from peculiar motions of rich clusters (Lauer \\& Postman 1992; Lauer et al. 1997) but cannot be used to explain the steep slope of near-infrared galaxy count. With this small local void introduced in our model, our underdensity fits yield voids that are at most a factor of 1.15 smaller in extent, but as much as 1.25 times more deficient in galaxies, than those listed in Tables 2 and 3. However, Wang, Spergel, \\& Turner (1998) used current knowledge of CMB anisotropies to show that a variation of a few percent between available measurements of $H_{0}$ and its true global value should be expected and that a variation as large as $10\\%$ would be possible for surveys with diameter $200h^{-1} Mpc$. For larger surveys scales, not much variation is expected. For example, for a survey with a $500h^{-1} Mpc$ radius, variations of at most $2\\%$ in expansion parameter and $13\\%$ in density are expected at the $95\\%$ confidence level. This would seem to mitigate against the large scale underdensity found by our fits if current cosmogonic theories are at least roughly valid. The goodness of fit ($Q > 0.01$) and maximum likelihoods values obtained with the second evolution model presented in section 3.3 seem to indicate that this is a viable alternative to the underdensity models. However, as we can see in Figure 4, our fit results show more evolution than can be accounted for by present evolution calculations. In fact, the (e+k)-correction from the fits is as much as 56\\% stronger than that of the 'Averaged' (e+k)-corrections for $\\Omega = 0.2$ used in our underdensity fits (for example, at $z = 0.5$, $\\Delta M_{(e+k)} = -1.12$ compared with $-0.74$) and 34\\% stronger than even the (e+k)-corrections for Sa galaxies alone (again for $\\Omega = 0.2$) which we took to represent an upper limit on the evolution (and for which $\\Delta M_{(e+k)} = -0.91$ at $z = 0.5$.) A more careful consideration of evolution as a source of the effect noticed by Huang97 and confirmed by our fits for all the data is necessary since our understanding of galaxy evolution is still incomplete. The recent publication of an extensive data base for Galaxy Evolution Modeling (Leitherer et al. 1996) might lead to larger (e+k)-corrections in the near-infrared. A model including both evolution and an underdensity might prove to be a more acceptable solution than either one alone, but cannot yet be studied with the data presently available. Redshift surveys should ultimately answer the question of whether or not a large region of the local universe is underdense. However, for the present, the K-band galaxy counts pose a significant puzzle: unless several independent determinations are giving similar but incorrect results, we must confront the possibility of either a cosmic density fluctuation of entirely unanticipated scale and amplitude or a serious deficiency in the best understood features of galaxy evolution and spectral energy distributions (or, of course, some combination of the two). Acknowledgments We are grateful to M. Fukugita and P. Garnavich for their help and useful comments. This research was supported by NSF grant AST94-19400. \\clearpage" + }, + "9802/astro-ph9802328_arXiv.txt": { + "abstract": "We present new photometric and spectroscopic data for the galaxies in the compact group known as Stephan's Quintet. We find the strongest evidence for dynamical perturbation in the spiral component NGC~7319. Most of the damage was apparently caused by nearby NGC~7320C which passed through the group a few $\\times$10$^8$ years ago. NGC~7318B is a spiral galaxy that shows evidence consistent with being in the early stages of a collision with the group. NGC~7317 and 18A are either elliptical galaxies or the stripped bulges of former spiral components. They show no evidence of past or present merger activity but are embedded in a luminous halo which suggests that they are interacting with the other members of the group. The low redshift galaxy NGC~7320 is most likely a late type, dwarf spiral projected along the same line of sight as the interacting quartet. ", + "introduction": "Stephan's Quintet (SQ: also known as Arp319 and VV228) is one of the most remarkable groupings of galaxies on the sky. It was the first compact galaxy group discovered and is certainly the most intensely studied object of that class. After every new observational effort unexpected aspects have emerged, adding to the puzzling nature of such apparently dense galaxy aggregate. SQ can be regarded as a prototype of the compact group class (it is H92 in the catalog of Hickson (1982). The puzzles connected with SQ began when its component galaxy redshifts were first measured (Burbidge \\& Burbidge 1961). First, one of the four higher redshift galaxies (NGC~7318B) was found to have a velocity almost 1000 \\kms~ lower than the other three raising questions about the groups dynamical stability. Then, NGC~7320 was found to show a redshift $\\sim$5700 \\kms~ lower than the mean of the other four. SQ became one of the key objects in the debate about the nature of the redshift (Sulentic 1983). It was suggested that NGC~7320 might be a physical companion of NGC~7331 a large, nearby Sb galaxy with similar redshift (\\cite{vdb61,a73}). Possible signs of this hypothesized interaction include a tidal tail that extends from NGC~7320 towards the SE and an HI deficiency noted by Sulentic and Arp (1983). At the same time Arp (1973) interpreted the tail as a sign of interaction between NGC~7320 and the higher redshift members of SQ. The question of the distance to the different galaxies in the group was specifically addressed by \\cite{bbgh73} and \\cite{sh74} who tried to use HI data to settle the question. The results turned out to be contradictory and later observations (\\cite{as80}) showed that the HI in SQ was displaced from the optical galaxies. A summary of the many papers dealing with distances to the SQ galaxies can be found in Sulentic (1983). This paper is concerned with the properties of the five galaxies that comprise SQ. We consider both their optical properties and the question of their normality, especially their past/present interaction state. The analysis of the properties of NGC~7318B are particularly relevant for the understanding the dynamical state of SQ because they suggest that it is a recent arrival in the system (see Moles, Sulentic \\& M\\'arquez 1997; MSM). The observations are presented in section 2 and analyzed galaxy-by-galaxy in section 3. The results are combined and summarized in section 4. ", + "conclusions": "" + }, + "9802/astro-ph9802058_arXiv.txt": { + "abstract": "We present Keck spectropolarimetry of the highly polarized radio-loud quasar 3CR\\,68.1 ($z=1.228, V=19$). The polarization increases from 5\\% in the red (4000\\AA\\ rest-frame) to $>$10\\% in the blue (1900\\AA\\ rest-frame). The broad emission lines are polarized the {\\em same} as the continuum, which shows that 3CR\\,68.1 is not a blazar as it has sometimes been regarded in the past. We also present measurements of the emission lines and a strong, blueshifted, associated absorption line system, as well as a detection at the emission-line redshift of \\ion{Ca}{2} K absorption, presumably from stars in the host galaxy. 3CR\\,68.1 belongs to an observationally rare class of highly polarized quasars that are neither blazars nor partially-obscured radio-quiet QSOs. Taking into account 3CR\\,68.1's other unusual properties, such as its extremely red spectral energy distribution and its extreme lobe dominance, we explain our spectropolarimetric results in terms of unified models. We argue that we have a dusty, highly inclined view of 3CR\\,68.1, with reddened scattered (polarized) quasar light diluted by even more dust-reddened quasar light reaching us directly from the nucleus. ", + "introduction": "3CR\\,68.1 ($z = 1.228, V = 19$) is a quasar with a unique combination of extreme properties. The optical-UV continuum is among the reddest of any known quasar, with $F_{\\nu} \\propto \\nu^{-6}$\\ in the UV (Boksenberg, Carswell, \\& Oke 1976; Smith \\& Spinrad 1980). 3CR\\,68.1 is one of the most highly polarized quasars, with a polarization in a broad bandpass from 3000-8000\\AA\\ of 7.5\\% $\\pm$ 1.3\\% at 52$\\arcdeg$ $\\pm$ 5$\\arcdeg$ (Moore \\& Stockman 1981, 1984). The lobes of 3CR\\,68.1 mark the quasar as a powerful radio source, but the extremely faint radio core ($\\approx 1$~mJy at 5\\,GHz, Bridle et al. 1994) makes this object one of the most lobe-dominant quasars. Under Unified Schemes, it is the most highly inclined 3CR quasar known (e.g., Orr \\& Browne 1982; Hoekstra, Barthel \\& Hes 1997). Extreme objects like 3CR\\,68.1 may provide clues as to how their properties may be related in less extreme AGNs. 3CR\\,68.1 has been classified as a blazar because of its high polarization and red color (e.g., Ledden \\& Odell 1985), but we thought this was unlikely because the degree of polarization and position angle did not change between measurements made 6 months apart (Moore \\& Stockman 1981, 1984), and the optical continuum level is not violently variable (Smith \\& Spinrad 1980). Also, it does not have the core-dominant radio structure always characteristic of radio-selected blazars. In fact, in the revised 3C sample (Laing, Riley, \\& Longair 1983; Hoekstra et al. 1997), it is the most lobe-dominant quasar with the largest projected lobe separation, large even among radio galaxies, suggesting a radio axis close to the sky plane. Most highly polarized QSOs (HPQs) can be classified as radio-quiet (including many BAL QSOs), or as radio-loud blazars. The polarization mechanism in radio-quiet QSOs is scattering (e.g., BAL QSOs: Hines \\& Wills 1995; Goodrich \\& Miller 1995; Cohen et al. 1995; Miller 1997; Ogle 1997; Schmidt et al. 1997; Brotherton et al. 1997; and IRAS non-BAL QSOs: Wills \\& Hines 1997; Wills et al. 1992), perhaps with a geometry similar to that found in polarized Seyfert 2 galaxies like NGC 1068 (e.g., Miller \\& Antonucci 1985, and the review by Antonucci 1993). In addition, many highly polarized, low-ionization BAL QSOs and IRAS QSOs appear reddened compared with UV-selected QSOs (Sprayberry \\& Foltz 1992). Do non-blazar radio-loud HPQs exist? OI 287 appears to be one -- a lobe-dominant radio-loud quasar, optically quiescent with a constant optical polarization of $\\approx$8\\% aligned parallel to the extended radio structure (Antonucci \\& Ulvestad 1988; Rudy \\& Schmidt 1988). The broad emission lines and continuum show similar degrees of polarization but the narrow emission lines are unpolarized (Goodrich \\& Miller 1988). Apparently OI 287 is a normal quasar seen only in scattered light because a thin, dusty disk obscures the central regions (Rudy \\& Schmidt 1988; Goodrich \\& Miller 1988; Antonucci, Kinney, \\& Hurt 1993). Certainly there are buried quasars in which broad lines and continuum are revealed only in scattered (polarized) light (e.g., Cygnus A, 3C 324, 3C 265, IRAS 09104+4109). Some of these are called radio galaxies. We hypothesized that the polarization mechanism in 3CR\\,68.1 is not optical synchrotron as in blazars, but rather scattering with dust absorption reddening the UV-optical spectrum. We therefore obtained spectropolarimetry of 3CR\\,68.1 with Keck II in order to test this hypothesis and determine the polarization mechanism. Furthermore, as part of the complete 3CR sample, whose selection is predominantly based on isotropic extended radio emission, an understanding of 3CR\\,68.1 can address the question of the prevalence of red quasars that might escape detection by usual optical selection criteria (e.g., Webster et al. 1995). \\S\\ 2 describes our observations and data reduction procedure. \\S\\ 3 presents our results and shows that the polarization increases dramatically to shorter wavelengths, that the broad emission lines are polarized at the same level as the continuum. We also discuss the absorption spectrum and our spectroscopic detection of the host galaxy of 3CR\\,68.1. \\S\\ 4 argues that scattering by dust or electrons is indeed the polarization mechanism, and that both the direct and scattered light are reddened. We show several models for the spectrum of 3CR\\,68.1, exploring the ranges in reddening, scattered light fraction, and host galaxy contribution. We briefly discuss how 3CR\\,68.1 fits into unification schemes. \\S\\ 5 is a summary. ", + "conclusions": "\\subsection{The Spectral Energy Distribution} 3CR\\,68.1 would not be found in an UV-excess survey. Some 10\\% of quasars in the 3CR sample, 3CR\\,68.1 being the most extreme, are significantly redder than bright optically selected quasars (Smith \\& Spinrad 1980). 3CR\\,68.1 has been observed in many wavebands. Figure 3 shows the spectral energy distribution (SED) of 3CR\\,68.1, from the far IR through the 2 keV X-rays, using measurements from the literature as noted in the caption. 3C 110, a normal `blue' lobe-dominated quasar, is also shown for comparison. The UV continuum is so weak that Bregman et al. (1985) argue that the large ratio of low-ionization to high-ionization line intensities is the result of photoionization by a separate soft X-ray component, giving rise to an extended partially ionized zone where strong low-ionization lines can be produced. A simpler explanation is that the quasar continuum and line spectrum are highly reddened, and emission-line clouds see a less obscured ionizing continuum than does an observer. The lack of a 2200\\AA\\ dust absorption feature is not a strong argument against reddening: no 2200\\AA\\ feature is seen in the Small Magellanic Cloud (SMC) reddening `law' (Prevot et al. 1984); even for Galactic dust, the 2200\\AA\\ feature can be filled in by scattering, in geometries possible for 3CR\\,68.1 (e.g., Witt, Thronson, \\& Capuano 1992). A 2200\\AA\\ dust feature is also absent in the spectra of BAL QSOs (Sprayberry \\& Foltz 1992; Hines \\& Wills 1995). Also, the composition of extragalactic dust may differ from that found in the Milky Way. Figure 4 compares our total flux spectrum with the photometry of Neugebauer et al. (1979) and Smith \\& Spinrad (1980). There is good agreement at blue wavelengths, and a ($\\approx$20\\%) difference at red wavelengths. Given the uncertanties in the flux calibration and our narrow slit, the general agreement is satisfactory. It seems probable that 3CR\\,68.1 has not varied significantly during the last twenty years. \\subsection{The Polarization Mechanism} The polarization mechanism in 3CR\\,68.1 is probably scattering by a non-spherical distribution of either small dust grains or electrons. Other mechanisms can be ruled out: \\begin{enumerate} \\item{Polarized synchrotron radiation: This mechanism fails to account for the broad-emission lines polarized the same as the continuum.} \\item{Galactic interstellar polarization: The interstellar polarization toward two stars close to the line of sight, SAO 55659 and SAO 55667 ($<$1$\\arcdeg$ from 3CR\\,68.1), is $\\leq$0.4\\% (D. Wills, private communication). } \\item{Host Galaxy interstellar polarization: The observed polarization shape and maximum differ significantly from a Serkowski Law, and as we show below, the maximum polarization is significantly greater than 9$\\times E(B-V)$, the polarization observed toward aligned dust-reddened lines of sight in the Milky Way.} \\end{enumerate} Scattering by electrons or small dust grains (2$\\pi a/\\lambda$ $<< 1$) can result in very high polarization (\\S\\ 4.3) that is essentially wavelength independent. Large grains, which polarize light with a stronger wavelength dependence, scatter less efficiently. We propose that the polarized spectrum arises from scattering of central quasar light, and that the rise in polarization to shorter wavelengths results from dilution by a redder spectrum. We further propose that the diluting spectrum is the quasar spectrum seen directly but reddened by dust, with some possible contamination by starlight from the host galaxy. It is likely that interstellar polarization from our galaxy or the host galaxy is small, so we do not include it. Hence the total observed spectrum of 3CR\\,68.1 can be described: \\begin{equation} Total_{\\lambda} = S_{\\lambda,A_S} + D_{\\lambda,A_D} + H_{\\lambda} + NLR, \\end{equation} \\noindent where $S_{\\lambda,A_S}$\\ is a scattered QSO spectrum reddened by an extinction $A_S$, $D_{\\lambda,A_D}$\\ is the direct QSO spectrum reddened by an extinction $A_D$, $H_{\\lambda}$\\ is the host galaxy spectrum, and $NLR$\\ is the spectrum of the narrow-line region which is unpolarized and is emitted probably from an extended region outside or coincident with the scattering region. The scattered spectrum is equal to the polarized spectrum divided by the polarization $p_s$, of the scattered light alone. To illustrate quantitatively the plausibility of our picture, and to constrain the fraction of scattered light and reddening towards the central quasar, we have matched composite model spectra to the total-light SED of 3CR\\,68.1. We first adjusted our spectrum to the photometry of Neugebauer et al. (1979). Free parameters are the intrinsic polarization of the scattered light, $p_s$, assumed to be wavelength-independent, and the fraction of host galaxy starlight. The strength of the \\ion{Ca}{2} K absorption feature suggests up to one quarter of the total light at 3900\\AA\\ rest-frame may be host galaxy starlight. We note that Palomar 5m image-tube plates show no extensions characteristic of a luminous host galaxy (Longair \\& Gunn 1975). We use the elliptical galaxy template of Kinney et al. (1996) to represent the starlight contribution. We represent the spectrum of the unreddened central quasar by a mean spectrum of lobe-dominated quasars (based on spectra of the HST sample described by Wills et al. 1995). The scattered light spectrum is derived from this mean spectrum by scaling and reddening it using an SMC-type extinction law (Prevot et al. 1984) with $A_V = 0.7$\\ to match the observed polarized flux density spectrum, then dividing by $p_s$. After subtracting the host galaxy starlight and scattered light from the SED, we are left with a spectrum that we fit with a reddened central quasar spectrum. Figure 5 presents four model composite spectra. Table 3 summarizes the model parameters, along with those of two intermediate cases. Models A, B, and C include no contribution from the host galaxy, and have $p_s = 80$\\%, 50\\%, and 20\\%. The direct line-of-sight reddening required is $A_V$ = 1.2, 1.3, and 1.6, respectively. Models A$_G$, B$_G$, and C$_G$ cover the same range in $p_s$, and include a host galaxy contribution of 25\\% of the total flux density at 3900\\AA. The deduced line-of-sight reddening is then $A_V$ = 1.1, 1.2, and 1.5, respectively. Also tabulated is the ratio of scattered to direct light after correction for reddening. Our comparison in Figure 5 extends over the range in our templates, from rest-frame $B$ to $H$; the fits over the range of our Keck data are in good agreement for the continuum, but if the SED from \\S\\ 4.1 is correct in the near-IR, Figure 5 shows that models without a large elliptical host galaxy contribution are preferred. Assuming H$_o$ = 50 km s$^{-1}$ Mpc$^{-1}$ and q$_o$ = 0, and using the observed spectrum and the known Galactic reddening, the absolute rest-frame $B$\\ magnitude for 3CR\\,68.1 is M($B$) = $-$26.6. Under the parameters of model A and dereddening the total light spectrum by $A_V = 1.2$, the intrinsic unreddened M($B$) = $-$28.2. This result is not very model dependent. Therefore 3CR\\,68.1 is one of the most optically luminous 3CR quasars, more in line with its radio luminosity, which is one of the highest among 3CR quasars. The radio and optical luminosities then follow the trend found by Wills \\& Brotherton (1995). \\subsubsection{Other Red, Scattered Light, Radio-Loud Quasars} There are few quasars with unusual combinations of properties similar to that of 3CR\\,68.1. As mentioned in \\S\\ 1, the quasar OI 287 (Antonucci, Kinney, \\& Hurt 1993; Goodrich \\& Miller 1988; Rudy \\& Schmidt 1988) has a high, stable optical polarization, a quiescent optical flux, blueshifted associated absorption, and is lobe dominant; OI 287 differs from 3CR\\,68.1 in that its polarization position angle is parallel to its radio axis. WN J0717+4611 (De Breuck et al. 1998) is another highly polarized red quasar, a steep spectrum radio source that is neither a blazar nor radio-quiet QSO. Its polarization is perpendicular to the radio jet axis. It has been proposed that these quasars are partially obscured, with red colors as a result of dust extinction, and polarization arising from scattering along less obscured lines of sight. 3CR\\,68.1 and objects such as these constitute a separate class of HPQ distinct from OVVs and radio-quiet QSOs. There are other radio-loud type 2 (NLR dominated) objects known, of quasar-like luminosities, with broad lines and continua seen in scattered light -- presumably `buried' quasars, e.g., Cygnus A (Ogle et al. 1997), 3C 324 (Cimatti et al. 1996), 3C 265 (Dey \\& Spinrad 1996), and IRAS 09104+4109 (Hines \\& Wills 1993). \\subsection{Geometry and Unified Schemes} 3CR\\,68.1 is a classic quasar, judging by its luminosity, and broad UV emission lines of typical equivalent width. At 178 MHz it is one of the three most powerful 3CR quasars in the compilation by Laing et al. (1983). The extreme dominance of the powerful lobes in 3CR\\,68.1, the largest projected lobe separation of any 3CR source, the reddest continuum of any quasar, and the highest scattered-light polarization of any 3CR source, immediately suggests a connection with highly inclined AGN in Unified Schemes. It is also of interest for testing Unified Schemes as it is a member of the only sample with complete optical identifications, selected by an essentially orientation-independent quantity directly related to the power of the central engine -- the low frequency (178 MHz) radio flux-density. Barthel (1989) based the arguments for the unification of radio-loud quasars and FR II radio galaxies on the 3CR sample, suggesting that radio galaxies are quasars at axial inclinations greater than the half opening angle of a geometrically and optically thick dusty torus, about 45$\\arcdeg$, so that their quasar nuclei are obscured. A variety of observational techniques have already shown that many 3CR radio galaxies harbor hidden quasar nuclei, and ionization and scattering `cones' have been found along the radio jet direction (e.g., Tran et al. 1998). In principle the polarization and reddening properties of the 3CR sample could be used to deduce the thickness of the obscuring torus as a function of inclination. Hill, Goodrich, \\& DePoy (1996) partially perform this exercise by looking for broad lines in the near infrared (where dust reddening is less) for a complete sample of 11 3C radio sources (8 narrow-line radio galaxies, two broad-line radio galaxies, plus 3C 273). They find that the majority of objects have their broad-line regions reddened by $A_V\\approx3.0$, and their narrow-line regions by $A_V\\approx1.5$. We investigate the properties of 3CR\\,68.1 in light of this Unification Scheme. Figure 6 compares the observational results with the results of a simple scattering model (after Brown \\& McLean 1977, see also Wills et al. 1992) in which we assume a uniformly filled cone of wavelength-independent scatterers and a scattering optical depth of unity appropriate for high polarization efficiency. The geometry of the cone is defined by the inclination $i$ of its axis to the observer's line-of-sight, and the cosine of its half-opening angle $\\theta_c$. The intrinsic polarization of light scattered towards the observer is $p_s$, and Scattered/Direct is the observed ratio of scattered to direct light assuming no obscuration. Grid lines of various $i$ and $cos \\theta_c$ are shown. From the position of the models of Table 3 on the grid, the cone half-opening angle is between $26\\arcdeg$ and $46\\arcdeg$, and the inclination $i > 41 \\arcdeg$. For an ideal `virtual' cone with an axis along the radio jets (the central engine axis) the polarization position angle (E-vector) is perpendicular to the radio axis, but the observed position angle for 3CR\\,68.1 is $\\approx60\\arcdeg$ to the radio axis. This discrepancy suggests a partially-filled virtual cone, or clumpy obscuration within a virtual cone. In this case the cone determined from Figure 6 will be defined by the scatterers that we `see'; it is tilted in the plane of the sky by 30$\\arcdeg$, and its opening angle is smaller than the size of the virtual cone. The smaller inclinations consistent with the polarization of 3CR\\,68.1 are near the Barthel model dividing value of 45$\\arcdeg$. The modest reddening towards the center (A$_V \\approx 1$ -- 1.5) is also between values for radio galaxies and UV-selected quasars. It may be helpful to realize that the observed continuum is especially steep because the UV is shifted to observed wavelengths, not because the reddening is especially high. The line-of-sight would be skimming the torus in this case. However, the extreme lobe-dominance and large lobe separation are not consistent with smaller inclinations. The larger inclinations, corresponding to jets almost in the plane of the sky, are consistent with the radio structure and the polarization results, but in this case we would expect the direct view of the center to be completely blocked by a thick torus. One way out is for the torus to be geometrically thin, a possibility also suggested for the other scattering-polarized radio-loud quasar OI\\,287. This might also tie in with the most luminous AGNs having the thinnest dusty disks (e.g., Antonucci et al. 1993). The strong, blueshifted, \\ion{Mg}{2} absorption in 3CR\\,68.1, identified by Aldcroft et al. (1993, 1994), and confirmed here, is sufficiently deep that the absorbing gas must lie at least in the direct path, and perhaps also in the scattered light path. In Unified Schemes lobe-dominant, high-inclination quasars are more likely to show absorption along the line-of-sight (for 3C quasars -- Wills et al. 1995, see also Aldcroft et al. 1993). 3CR\\,68.1 follows this trend. Several modifications of the simple Unified Model for radio galaxies and quasars, involving the effects of absorption towards the central source, the radio core, the scatterers and the NLR, are suggested in the case of 3CR\\,68.1: (i) even for highly inclined AGNs like the radio galaxies, the radio core dominance is among the smallest (log\\,R = $-3.2$). In the lower inclination case, a possibility is that the reddened line-of-sight combined with the very weak radio core is the result of radio free-free absorption by dusty, ionized gas. This could be tested by determining the shape of the radio spectrum, and also by searching for X-ray absorption edges of O\\,VII and O\\,VIII (as for IRAS\\,13349+2438 -- Beichman et al. 1986; Brandt et al. 1997; see also Leighly et al. 1997; Grupe et al. 1997). (ii) the polarization position angle is not perpendicular to the radio jets as discussed above, suggesting a partially filled or partially obscured cone, and (iii) the narrow-line (NLR) equivalent widths are similar to those in our mean unreddened quasar spectrum, and also similar to the composite lobe-dominated quasar spectrum of Baker \\& Hunstead (1995) formed from the Molonglo sample. If 3CR\\,68.1 were reddened by the amounts implied by our modeling, then the continuum is suppressed by 20 to 30 times. If the low-ionization NLR extends well beyond the dusty torus, the NLR should suffer much less reddening and EW(NLR) should be enhanced by 20 to 30 times. The reddening of the NLR in the Hill et al. radio-galaxies, the reddened scattered spectrum of 3C\\,68.1, and a partially obscured cone, all suggest significant absorption. We reach a limit to this simple interpretation if there is significant forward-scattered, less-polarized light; in this case the obscuration we derive from the steepening of the UV-optical spectrum would be an underestimate. But then the geometry of the simple Unified Scheme breaks down. In this regard, 3CR\\,68.1 makes an especially interesting and important comparison with the (only) other scattered-light, luminous radio-loud quasar OI\\,287. In OI\\,287 the scattered light is thought to be scattered by electrons in the outer regions of a geometrically thin, dusty disk that obscures the center completely at UV-optical wavelengths. This explains its polarization parallel to the radio jets. 3CR\\,68.1 also shows polarization that is not exactly perpendicular to the radio jets; a polarized component as in OI 287 provides an alternative explanation to a partially filled or partially obscured scattering cone. Overall 3CR\\,68.1 fits well into the Unified Scheme, although it falls short of being a textbook example in several potentially interesting respects. These `shortcomings' may in the future provide opportunities to probe intrinsic obscuring structures common to all quasars." + }, + "9802/astro-ph9802091_arXiv.txt": { + "abstract": "Recent determinations of the upper mass limit to the local initial mass function (IMF) claim a value of $m_{\\rm U}=50\\pm 10$ M$_\\odot$, based upon direct comparisons of the observed oxygen and iron abundances in metal-poor stars with the predicted stellar yields from Type II supernovae (SNe). An unappreciated uncertainty in these analyses is the input physics intrinsic to each SNe grid, and its effect upon stellar nucleosynthesis. We demonstrate how such uncertainties, coupled with the uncertain metal-poor halo star normalization, while allowing us to set a \\it lower \\rm bound to $m_{\\rm U}$ of $\\sim 40$ M$_\\odot$, nullifies any attempt at constraining the \\it upper \\rm bound. ", + "introduction": "\\label{introduction} The initial mass function (IMF), a measure of the distribution of masses, at formation, of a given stellar generation, is one of the key components of galaxy evolution modeling. Despite its importance, accurate determination of the IMF remains one of the most elusive problems in modern astronomy. This elusiveness is manifested both in ongoing attempts to understand its physical underpinings (\\eg Padoan, Nordlund \\& Jones 1997), as well as simply characterizing its shape and mass limits from an observational tack (\\eg Kroupa, Tout \\& Gilmore 1993; Scalo 1986). In its simplest form, the IMF can be considered a power law of the form $m\\phi(m){\\rm d}m\\propto m^{-(1+x)}{\\rm d}m$, where $m\\phi(m){\\rm d}m$ is the number of stars born in the mass interval $m\\rightarrow m+{\\rm d}m$. The goal for theorists and observers, alike, is then the determination of the slope of this function $x$, as well as its upper and lower limits ($m_{\\rm U}$ and $m_\\ell$, respectively). While the slope of the IMF, at least in the solar neighborhood (and for masses greater than a few solar masses), would appear to lie somewhere in the range $x\\approx 1.3$ (Salpeter 1955) to $x\\approx 1.7$ (Kroupa, Tout \\& Gilmore 1993), and the lower mass limit is close to $m_\\ell\\approx 0.2$ M$_\\odot$ (Bahcall \\etal 1994), the upper mass limit $m_{\\rm U}$ still remains \\it highly \\rm uncertain. Even a cursory examination of the literature corroborates this point, with values in the range $m_{\\rm U}\\approx 20\\rightarrow 200$ M$_\\odot$ suggested by a variety of direct and indirect techniques (\\eg Maeder \\& Meynet 1989; Klapp \\& Corona-Galinda 1990; Pagel \\etal 1992; Maeder 1992; Massey \\etal 1995; Kudritzki 1997). A hybrid approach to determining $m_{\\rm U}$, combining predictions of the \\it theoretical \\rm yields from Type II supernovae (SNe), with the \\it observed \\rm abundances in the old metal-poor stars (\\ie those which bear the clear imprint of yield ``pollution'' from these same SNe, with no ``dilution'' from Type Ia SNe, whose progenitor lifetimes are considerably longer than the Type II timescales), has been the subject a recent series of papers (Tsujimoto \\etal 1995,1997; Yoshii, Tsujimoto \\& Nomoto 1996). The premise here is that because Type II SNe $\\alpha$-element (\\eg O, Mg, Ne) yields are a strong function of progenitor mass, whereas products of explosive burning (\\eg Fe, Si, Ca) are less so, the IMF-weighted average of their ratios must necessarily also depend strongly upon $x$ and $m_{\\rm U}$. Tying these yield ``averages'' to the halo abundances then, in principle, provides a unique indirect probe on the upper mass limit to the IMF.\\footnote{Tsujimoto \\etal (1997) extend this ``hybrid'' approach to simultaneously constrain the IMF mass limits $m_{\\rm U}$ and $m_\\ell$, as well as the slope $x$. We shall only be concerned with the $m_{\\rm U}$ determination in what follows, primarily for brevity, but also because the lower mass constraint rest squarely upon uncertain photometric calibrations.} Following this technique, Tsujimoto \\etal (1997) recently concluded that the upper mass limit to the IMF in the solar neighborhood is $m_{\\rm U}=50\\pm 10$ M$_\\odot$. To do so, they made explicit use of the Tsujimoto \\etal (1995) compilation of Type II SNe yields. What was not fully appreciated in their study though was just how dependent their result was to this \\it particular \\rm yield compilation and the adopted halo abundance normalization. It is to this lack of appreciation that our current study is addressed. After providing a minimal introduction to the model ingredients in Section \\ref{analysis1}, we demonstrate in Sections \\ref{analysis2} and \\ref{analysis3} that this technique results in $m_{\\rm U}=50\\pm 10$ M$_\\odot$ \\it only \\rm for the Tsujimoto \\etal (1995) yields, combined with a halo normalization of [O/Fe]$_{\\rm h}=+0.41$. Duplicating the analysis with ``competing'' yield compilations which sample a wide variety of convection and mass-loss treatments (\\eg Woosley \\& Weaver 1995; Langer \\& Henkel 1995; Arnett 1996), clearly demonstrates that Tsujimoto \\etal (1997) have significantly underestimated the uncertainty associated with thir determination of $m_{\\rm U}$. Our results are summarized in Section \\ref{summary}. ", + "conclusions": "\\label{summary} T97 have recently revitalized interest in using IMF-weighted Type II SNe yields as a direct probe of said IMF's upper mass limit $m_{\\rm U}$, by comparison with the observed abundance ratios in metal-poor Galactic stars. The beauty of this technique lies partially in its simplicity -- for a given IMF slope, there is effectively only one free parameter -- the yield source. Adopting the T95 yields, T97 found $m_{\\rm U}=40\\rightarrow 60$ M$_\\odot$, for reasonable IMF slopes (\\ie $x=1.3\\rightarrow 1.6$). The primary concern we have regarding T97's analysis lies in their underlying assumption that \\it all \\rm of the uncertainty in the stellar model physics could be encapsulated in a 30\\% (0\\%) error budget for oxygen (iron). It should be obvious from Figures \\ref{fig:fig1} and \\ref{fig:fig2}, and Langer (1997), for example, that this assumption is incorrect, and that a more realistic error budget would allow for up to an order of magnitude greater leeway. Stellar models are simply not developed to the level that is inherently assumed by T97 -- convection, overshooting, mass-loss, reaction rates, metallicity, C/O-core masses, fallback onto the remnant -- each conspire to increase the uncertainties to the degree reflected by the yields shown in Figures \\ref{fig:fig1} and \\ref{fig:fig2}. A secondary concern is T97's inherent assumption that the halo normalization [O/Fe]$_{\\rm h}=+0.41$ has no associated uncertainty. Since values as high as [O/Fe]=+0.6 are still favored by some, this $\\sim 60$\\% uncertainty should be taken into account. For a Salpeter (1955) slope, there is (roughly) a one-to-one correspondence between the halo normalization uncertainty and the corresponding predicted upper mass limit uncertainty. While we \\it do \\rm agree with T97 that the \\it lower \\rm limit to $m_{\\rm U}$ is $\\sim 40$ M$_\\odot$ (or $\\sim 30$ M$_\\odot$, if we adopt the extreme A78 yields), our more realistic exploration of input physics ``space'' demonstrates that we simply cannot constrain the \\it upper \\rm limit to any useful accuracy. Taken together, we can only conclude that, by this technique alone, $m_{\\rm U}\\simgt 40$ M$_\\odot$, for IMF slopes $x=1.3\\rightarrow 1.6$. Fixing the IMF slope to that of Salpeter (1955), we can only constrain $m_{\\rm U}$ to lie somewhere between $\\sim 40$ M$_\\odot$ and $\\sim 140$ M$_\\odot$. Finally, it would appear to be difficult to reconcile any $m_{\\rm U}\\simlt 100$ M$_\\odot$ with either the A96$_{\\rm ZAMS}$ or WW95;+0.60 halo normaliztion grids. While promising (provided existing discrepancies in Type II SNe yields are eliminated), at the present time, unfortunately, this technique, by itself, does not substantially improve or constrain our understanding of the upper mass limit to the solar neighborhood IMF." + }, + "9802/hep-ph9802271_arXiv.txt": { + "abstract": "We examine the $X$-ray spectrum from the decay of the dark-matter moduli with mass $\\sim {\\cal O}(100)$keV, in particular, paying attention to the line spectrum from the moduli trapped in the halo of our galaxy. It is found that with the energy resolution of the current experiments ($\\sim 10$\\%) the line intensity is about twice stronger than that of the continuum spectrum from the moduli that spread in the whole universe. Therefore, in the future experiments with higher energy resolutions it may be possible to detect such line photons. We also investigate the $\\gamma$-ray spectrum emitted from the decay of the multi-GeV moduli. It is shown that the emitted photons may form MeV-bump in the $\\gamma$-ray spectrum. We also find that if the modulus mass is of the order of 10 GeV, the emitted photons at the peak of the continuum spectrum loses their energy by the scattering and the shape of the spectrum is significantly changed, which makes the constraint weaker than that obtained in the previous works. ", + "introduction": "Superstring theories\\cite{Green}, which may be the most attractive candidates to unify all known interactions including gravity, have a number of flat directions, called moduli fields, in a large class of classical ground states\\cite{Green}. These moduli fields $\\phi$ continuously connect infinitely degenerate supersymmetric vacua and they are generally expected to acquire their masses $m_\\phi$ of the order of the gravitino mass $m_{3/2}$ once supersymmetry breaking effects are included\\cite{Carlos-Casas-Quevedo-Roulet}. These moduli fields cause different kinds of cosmological problems \\cite{Coughlan,K-Y} depending on values of their masses. At present the thermal inflation proposed by Lyth and Stewart \\cite{Lyth-Stewart} seems to be the most plausible solution to the problems. In recent articles\\cite{H-K-Y,A-H-K-Y}, we have shown by postulating the thermal inflation that only two regions of the moduli masses, $m_\\phi \\lesssim $ 500 keV and $m_\\phi \\gtrsim {\\cal O}$(100) GeV, are cosmologically viable. In particular, the lighter mass region is more interesting since the original Affleck-Dine baryogenesis \\cite{Affleck-Dine} does work here as shown first by de Gauv\\^{e}a, Moroi and Murayama\\cite{G-M-M}. On the contrary, for $m_\\phi \\gtrsim {\\cal O}$(100) GeV we must invoke a variant type of Affleck-Dine baryogenesis\\cite{S-K-Y} which has not been, however, well investigated yet. If the moduli masses lie indeed in the region $m_\\phi \\simeq 10^{-2}$ keV--200 keV there is an intriguing possibility\\cite{A-H-K-Y} that the moduli fields are the dark matter in our universe. Since the thermal inflation produces a tremendous amount of entropy at the late epoch of the universe's evolution to dilute the moduli density substantially, there seems to be no candidate left for the dark matter beside the moduli themselves \\footnote{ The axion with high values of decay constant $F_a \\simeq 10^{15}$--$10^{16}$ GeV could be another candidate for the dark matter\\cite{L-S-S-S}. }. This would encourage us to consider the hypothesis of moduli being the dark matter in the universe. In this paper we calculate spectrum of background X-rays emitted from the moduli dark matter and find that the spectrum is constituted of two distinct parts: one comes from the cosmic moduli filling homogeneously the whole universe and the other from the moduli condensed on the dark halo in our galaxy. The former has a relatively broad spectrum due to the redshift effect and the latter has a peak in the energy spectrum. We show that the peak in the X-ray spectrum can be detectable in future experiments if the moduli masses $m_\\phi$ are around 100 keV. We also briefly comment on $\\gamma$-ray spectrum emitted from more massive moduli of $m_\\phi \\simeq$ 1 -- 10GeV, since this multi-GeV mass region is marginally allowed \\cite{H-K-Y,A-H-K-Y} if one assumes somewhat smaller values of the initial amplitudes of moduli fields, $\\phi_0 \\simeq$ (0.01 -- 0.1)$M_G$, where $M_G$ is the gravitational scale $M_G \\simeq 2.4 \\times 10^{18}$ GeV. We find that the $\\gamma$-rays emitted from such moduli make a large bump in multi-MeV region and the shape of the spectrum depends heavily on the masses of moduli. ", + "conclusions": "In this paper we have examined the photon spectra from the decay of the cosmic modulus field. First we have considered the modulus mass region $m_{\\phi} \\simeq$ $10^{-2}$keV--200 keV. This region is interesting because the modulus field can be the dark matter in our universe. We have calculated the $X$-ray continuum spectrum from the decay of the dark-matter moduli that spread homogeneously in the whole universe and the line spectrum from the dark-matter moduli trapped in the halo of our galaxy. It is found that with the energy resolution of the current experiments ($\\sim 10$\\%) the line intensity is about twice stronger than that of the continuum spectrum in the wide region of the sky. If the modulus mass is around 100 keV, both intensities are comparable with the present observed photon backgrounds. Therefore, in the future experiments with higher energy resolutions it may be possible to detect the line photons produced by the decay of dark-matter moduli in our halo. Moreover, by measuring the dependence of the line intensity on the galactic longitude and latitude, we will be able to confirm the origin, i.e. it comes from the halo of our galaxy rather than from the whole universe. We have also investigated the $\\gamma$-ray spectrum emitted from the decay of the multi-GeV modulus field. In this modulus mass region, the emitted photons are redshifted and have a peak in the MeV region of the spectrum. Thus we may observed those photons as a MeV-bump in the $\\gamma$-ray backgrounds. The produced high energy photon may be scattered off the background photons and lose their energy. It is found that the effect of the scattering is negligible for modulus with mass less than ${\\cal O}(1)$GeV. However, if the modulus mass is of the order of 10 GeV, the emitted photons at the peak of the continuum spectrum loses their energy by the scattering and the shape of the spectrum is significantly changed. This makes the constraint from the present observed $\\gamma$-ray backgrounds weaker than the result in Ref.\\cite{A-H-K-Y}." + }, + "9802/astro-ph9802085_arXiv.txt": { + "abstract": "A multivariate analysis of gamma-ray burst (GRB) bulk properties is presented to discriminate between distinct classes of GRBs. Several variables representing burst duration, fluence and spectral hardness are considered. Two multivariate clustering procedures are used on a sample of 797 bursts from the Third BATSE Catalog: a nonparametric average linkage hierarchical agglomerative clustering procedure validated with Wilks' $\\Lambda^*$ and other MANOVA tests; and a parametric maximum likelihood model-based clustering procedure assuming multinormal populations calculated with the EM Algorithm and validated with the Bayesian Information Criterion. The two methods yield very similar results. The BATSE GRB population consists of three classes with the following Duration/Fluence/Spectrum bulk properties: Class I with long/bright/intermediate bursts, Class II with short/hard/faint bursts, and Class III with intermediate/intermediate/soft bursts. One outlier with poor data is also present. Classes I and II correspond to those reported by Kouveliotou et al. (1993), but Class III is clearly defined here for the first time. ", + "introduction": "As very few gamma-ray burst (GRB) sources have astronomical counterparts at other wavebands, empirical studies of GRBs have been largely restricted to the analysis of their gamma ray properties: bulk properties such as fluence and spectral hardness, and evolution of these properties within a burst event (Fishman \\& Meegan 1995). While bursts exhibit a vast range of complex temporal behaviors, their bulk properties appear simpler and amenable to straightforward statistical analyses. Studies fall into two categories: examination of whether GRB bulk properties comprise a homogeneous population or are divided into distinct classes; and search for relationships between bulk properties. Both types of study may lead to astrophysical insight, just as the distinction between main sequence stars and red giants and the measurement of a luminosity-mass relation along the main sequence assisted the development of stellar astrophysics early in the century. The most widely accepted taxonomy of GRBs is the division between short-hard and long-soft bursts proposed by Dezalay et al. (1992) and Kouveliotou et al. (1993, henceforth K93). K93 noticed a bimodality in the burst duration variable $T_{90}$ (time within which 90\\% of the flux arrived), suggesting the presence of two distinct types of bursts separated at $T_{90} \\simeq 2$ sec. The short bursts have systematically harder gamma-ray spectra than longer bursts. The two groups seemed indistinguishable in most other bulk properties, although the larger group of long-soft bursts may have a subclass with a different fluence distribution (i.e., different $$; Katz \\& Canel 1996) and the groups may have different Galactic latitude distributions (Belli 1997). Other researchers point to small groups of bursts with distinctive properties such as the soft-gamma repeaters (Norris et al. 1991), two possible classes with differing short-timescale variability (Lamb, Graziani \\& Smith 1993), fast-rise exponential-decay bursts (Bhat et al. 1994), and two types of bursts with different ratios of total fluence and $>$300 keV fluence (Pendleton et al. 1997). A variety of relationships between burst properties have also been reported. Norris et al. (1995) find an anti-correlation between $T_{90}$ (calculated after wavelet thresholding) and peak intensity, consistent with a cosmological time dilation. However, a positive correlation between $T_{90}$ and total fluence is also seen which does not agree with the simplest cosmological interpretation (Lee \\& Petrosian 1997). Additional reported relationships include: $T_{90}$ correlated with peak heights (Lestrade 1994), peak energy correlated with peak flux (Mallozzi et al. 1995), and peak duration anticorrelated with gamma-ray energy (Fenimore et al. 1995). Most of these studies suffer from a failure to treat all of the bulk property variables in an unbiased and quantitative way. Astronomers typically examine univariate or bivariate distributions, sometimes constructing composite variables (such as hardness ratios) with pre-determined relationships to include one or two additional variables. But it is quite possible that the complex astrophysics producing GRBs will not manifest themselves in simple bivariate plots, just as the division between short-hard and long-soft bursts is not evident in spectral variables alone (Pendleton et al. 1994). GRB catalogs, like most multiwavelength astronomical catalogs, are multivariate databases and should be treated with multivariate statistical methods that can objectively and effectively uncover structure involving many variables (Feigelson \\& Babu 1997). Two previous studies take a fully multivariate approach to understanding GRB bulk properties. Baumgart (1994) constructs a neural network taxonomy of 99 GRBs from the PVO satellite using 26 variables representing both bulk burst properties and detailed temporal characteristics (e.g. number of peaks, fractal dimension, wavelet transform crossings) and finds two or three distinct GRB classes. Bagoly et al. (1997) perform principal components and factor analyses of nine bulk property variables using 625 GRBs from the BATSE 3B catalog. They find that the relationships in the database are determined principally by only three variables: an appropriately weighted fluence, a weighted burst duration, and (to a lesser extent) flux in the highest energy bin. We note, however, that it can be dangerous to look for correlations prior to classifying (or establishing the homogeneity of) the population. While the anticorrelation between hardness ratio and burst duration seen in full samples (K93) may be the manifestation of a single astrophysical process, it may alternatively reflect differences between distinct processes. The latter possibility is suggested by a reported hardness-duration positive correlation within the long-soft class of bursts (Dezalay et al. 1996; Horack \\& Hakkila 1997). Most multivariate analyses thus begin with a study of homogeneity and classification, and then investigate the variance-covariance structure (i.e. correlations) within each class. This paper describes a multivariate analysis of GRBs from the Third BATSE Catalog (Meegan et al. 1996). After defining the sample (\\S 2), we start with a simple statistical description of the variables and their bivariate relationships for the entire dataset (\\S 3). We then seek distinct types of clusters in two ways. First, a standard nonparametric agglomerative hierarchical clustering analysis is performed (\\S 4) which reveals three distinct classes. The statistical significance of the third cluster is validated, under Gaussian assumptions, with MANOVA tests. Second, a parametric maximum likelihood model-based clustering procedure is adopted which reveals the same three groups and indicates strong evidence for the presence for three rather than two groups (\\S 5). The variance-covariance structure of each group is then examined (\\S 6). Results are synthesized in the discussion (\\S 7). Throughout the paper, we discuss our mathematical techniques to help the reader understand the complexities of multivariate analysis. From the vast literature in this subject, we recommend the following monographs for interested readers: Johnson \\& Wichern (1992) and Jobson (1992) for overviews of applied multivariate analysis; Hartigan (1975), Jain \\& Dubes (1988) and Kaufman \\& Rousseeuw (1990) for multivariate clustering algorithms; Murtagh \\& Heck (1987) and, more briefly, Babu \\& Feigelson (1996) and Feigelson \\& Babu (1997), for multivariate methodology in astronomy. ", + "conclusions": "We thus find, using clustering and validation methods with different mathematical underpinnings, that three classes of GRBs are present in our large subset of the Third BATSE Catalog. Most of the structure can be found using three fundamental burst properties, Duration/Fluence/Spectrum. The class properties and relation to previous research can be briefly summarized as follows: \\begin{description} \\item[Class I] These long/bright/intermediate bursts correspond to the well-known populous long-soft class of K93 and others. Within this group, we do not confirm a hardness-duration correlation reported by Dezalay et al. (1996) and Horack \\& Hakkila (1997). \\item[Class II] This short/faint/hard group corresponds to the short-hard burst type of K93 and others. Fluence-duration and fluence-hardness correlations may tentatively be present within the class. Note that while the mean location of this type is consistent in the two clustering schemes, its size and population (e.g. 1/2 or 1/4 that of Class I) differs between clustering algorithms. \\item[Class III] The discovery of this group with intermediate/intermediate/soft properties is the principal result of this study. The group is easily distinguished in the projections of Figure 3, but can also be discerned in some panels of Figure 1. For example, it lies between Classes I and II in the $T_{50}-H_{32}$, $T_{90}-F_{tot}$ and $T_{90}-H_{321}$ scatter plots. In the univariate $T_{90}$ distribution shown in Figure 5, Class III accounts for most, but not all, of the bursts in the small peak around $2 < T_{90}<5$ sec between the major short and long duration peaks. It is possible that our Class III is related to the class of no-high-energy (NHE) burst and peaks discussed by Pendleton et al. (1997). These bursts have unusually weak $F_4$ emission, soft $50-300$ keV spectra, and low $F_{tot}$. However, the NHE class does not appear to exhibit a clear duration segregation from other bursts as we find for Class III. Class III does not appear to be the third cluster found by Baumgart (1994, see his Table 3), but the high dimensionality of his analysis prevents a simple comparison with our low dimensionality study. \\item[Outlier] BATSE trigger event 2757, burst 3B 940114, is the outlier in the nonparametric analysis of \\S 4 and is clearly visible in many projections in Figures 1 and 3.\\footnote{The model-based analysis of \\S 5 cannot locate clusters with very few members and assigned this event to Class II. An extension of model-based clustering that models outliers as Poisson noise can do this (see Banfield and Raftery 1993, Dasgupta and Raftery 1998), but it does not seem necessary in this application.} It has an exceedingly soft hardness ratio and short burst duration. But examination of the original BATSE database shows that the $F_1-F_4$ fluxes are very weak with large measurement uncertainties. The published 3B catalog gives only an upper limit to its total fluence and no estimate of its hardness ratio (Meegan et al. 1996). The unusual properties of this burst are thus illusory and are due to its very weak fluence. \\end{description} The multivariate analysis described here is not comprehensive and may not have uncovered all of the structure in the Third BATSE Catalog of bulk GRB properties. Our reduction of dimensionality may have been too severe, omitting, for example, the potentially important $F_4$ as a distinct variable (Pendleton et al. 1997; Bagoly et al. 1997). Many methodological options were not exercised. For example, it would be valuable to repeatedly apply the $k$-means partitioning algorithm to the database under the assumption that three clusters are present (see Murtagh 1992 for an astronomical application of this method), check for skewness or kurtosis in the clusters, and undertake an oblique decision tree analysis to give analytical formulation to hyperplanes separating the clusters (see White 1997). Codes for these and many other multivariate techniques are publicly available through the Web metasite {\\it StatCodes} at www.astro.psu.edu/statcodes. However, the efforts described here are far more capable of finding and quantifying clustering in the database than most previous analyses (\\S 1). Previous studies have been based on qualitative rather than quantitative procedures for identifying structures, and provide no statistical validation of their claims. It is thus not surprising that we uncovered structure missed by previous researchers. In particular, our confidence in the presence of a third cluster, Class III, is strong. Two completely independent mathematical procedures (\\S 4 and \\S 5) found very similar structure, each validated with high statistical confidence. It is possible that the clustering reported here is indeed present in the database, but does not have an astrophysical origin. The complex triggering mechanism of the BATSE instrument mechanism and biases in bulk property values at low signal-to-noise ratio are two problems that probably affect the multivariate structure. We have investigated one manifestation of the latter effect using $T_{90}^d$ and found no effect on our results. Instrumental biases generally affect the number of bursts found in some regions of the multivariate hyperspace (thereby biasing log $N$ - log $S$ distributions) and may alter the location of clusters, but are unlikely to cause the appearance of clustering that is not present in the underlying population. Nonetheless, since the BATSE instrument identifies bursts on three separate timescales, it is possible that the third cluster here is related to a selection effect associated with the BATSE triggering mechanisms. We conclude that the Third BATSE Catalog shows three statistically significant types of bursts (Duration/Fluence/Spectrum): Class I GRBs are long/bright/intermediate, Class II GRBs are short/faint/hard, and Class III GRBs are intermediate/intermediate/soft. Unless the separation of Class III from the other types is due to some subtle BATSE instrumental effect, these types are likely to be real and their existence should be considered an important input into astrophysical theories for GRBs. For example, the three types may reflect different types of external environments and internal shocks in relativistic fireball models (M\\'esz\\'aros \\& Rees 1993; Panaitescu \\& M\\'esz\\'aros 1998). Note that statistical anlaysis is unable to determine whether burst types represent fundamentally different astrophysical processes or distinct conditions within a single astrophysical model. Our results can be confirmed and extended in two fashions. First, the analysis described here can be validated with several hundred more bursts collected by BATSE since the September 1994 cutoff in the database used here. Second, following Baumgart (1994), the dimensionality of the problem can be enlarged to include detailed characteristics of the burst temporal behaviors. Burst smoothness $vs.$ peakiness, characteristic wavelet scales, spectral evolution, and other parameters can be included. With this enlarged database, one can perform both an unsupervised exploratory cluster analysis similar to that described here, and MANOVA-type analyses that assume the existence of the three groups to determine whether the clusters have distinctive temporal properties." + }, + "9802/astro-ph9802170_arXiv.txt": { + "abstract": "We present spatially-resolved, moderate-resolution spectrophotometry of the recurrent nova T~Pyx and a portion of the surrounding shell. The spectrum extracted from a strip of width $10''$ centered on the star shows well-known, strong emission lines typical of old novae, plus a prominent, unfamiliar emission line at $\\lambda$6590. This line, and a weaker companion at $\\lambda$6540 which we also detect, have been previously reported by Shahbaz et al., and attributed to Doppler-shifted H$\\alpha$ emission from a collimated jet emerging from T~Pyx. We demonstrate that these lines are instead due to [N\\,II] $\\lambda\\lambda$6548, 6584 from a complex velocity field in the surrounding nebula. The comments of past workers concerning the great strength of He\\,II $\\lambda$4686 in T~Pyx itself are also reiterated. ", + "introduction": "The recurrent nova T~Pyxidis was discovered by Leavitt (1914); a history of the object, and some recent observations, may be found in Webbink et al. (1987) and Shara et al. (1989). A compact ($\\sim10''$) nebula surrounding the star is described by Duerbeck \\& Seitter (1979) and Williams (1982), and a larger, fainter shell is observed by Shara et al. (1989). {\\it Hubble Space Telescope} imagery of the T~Pyx nebula (Shara et al. 1997) reveals an exceptionally complex, clumpy structure on subarcsecond spatial scales, and the brightness of at least some of these knots varies on timescales of months. As pointed out by these authors, the {\\it HST} data vividly demonstrate that ``shell\" is a quite misleading term for the extended structure near T~Pyx, which in fact consists of literally thousands of discrete lumps, most bright in [N\\,II] emission. Recently an interesting spectrum of T~Pyx has been presented by Shahbaz et al. (1997). They call attention to a strong, unfamiliar emission line at $\\lambda6593$, and a weaker feature at $\\lambda6539$. They interpret these features as Doppler-shifted H$\\alpha$ lines from a collimated jet emerging from T~Pyx, at velocities of +1400 and $-1100$~km~s$^{-1}$. As there is no {\\it a priori} reason to expect these particular velocities, and no second spectral feature to confirm them, care must obviously be taken when accepting this interpretation; virtually any unidentified emission line, regardless of wavelength, can be attributed to such a model. Nonetheless this interpretation if correct is very exciting: as stressed by Shahbaz et al. (1997), this would make T~Pyx the first short-period cataclysmic variable with a jet. Livio (1998), who comments that ``jet lines have now been observed unambiguously in the recurrent nova T~Pyx,\" stresses that there would be profound implications on models of jet formation. In an effort to clarify this unprecedented interpretation, we obtained further spectra of T~Pyx and the surrounding nebula. Although we verify the existence of the unusual emission lines, the features unfortunately prove to have a less exotic origin than that suggested by Shahbaz et al. ", + "conclusions": "We have presented spatially-resolved spectrophotometry of T~Pyx and portions of the surrounding nebula. Recent {\\it HST} images have vividly stressed how spatially complex the extended structure is. From our data, it is also clear that the shell is spectrally complex. We believe that there is little evidence that the prominent $\\lambda$6590 emission line and its weaker companion at $\\lambda$6540 are due to H$\\alpha$ from a collimated jet ejected by the star (Shahbaz et al. 1997). Rather, we argue, based on the detection of multiple emission lines from different species at consistent velocities, that these lines are instead [N\\,II] from a few, or possible even many, discrete knots. The large He\\,II $\\lambda$4686/H$\\beta$ ratio, probably due to T~Pyx itself, also bears further scrutiny." + }, + "9802/astro-ph9802346_arXiv.txt": { + "abstract": "Using the Long Wavelength Spectrometer (LWS) onboard the Infrared Space Observatory (ISO), we have observed thermal water vapor emission from a roughly circular field of view approximately 75 arc seconds in diameter centered on the Orion BN-KL region. The Fabry-Perot line strengths, line widths, and spectral line shifts observed in eight transitions between 71 and 125$\\mu$m show good agreement with models of thermal emission arising from a molecular cloud subjected to a magnetohydrodynamic C-type shock. Both the breadth and the relative strengths of the observed lines argue for emission from a shock rather than from warm quiescent gas in the Orion core. Though one of the eight transitions appears anomalously strong, and may be subject to the effects of radiative pumping, the other seven indicate an H$_2$O/H$_2$ abundance ratio of order $5\\times 10^{-4}$, and a corresponding gas-phase oxygen-to-hydrogen abundance ratio of order $4\\times 10^{-4}$. Given current estimates of the interstellar, gas-phase, oxygen and carbon abundances in the solar vicinity, this value is consistent with theoretical shock models that predict the conversion into water of all the gas-phase oxygen that is not bound as CO. The overall cooling provided by rotational transitions of H$_2$O in this region appears to be comparable to the cooling through rotational lines of CO, but is an order of magnitude lower than cooling through H$_2$ emission. However, the model that best fits our observations shows cooling by H$_2$O and CO dominant in that portion of the post-shock region where temperatures are below $\\sim 800$\\,K and neither vibrational nor rotational radiative cooling by H$_2$ is appreciable. ", + "introduction": "On October 6, 1997, we observed the Orion BN-KL region with ISO (cf Kessler {\\it et al.}, 1996) from 06:02:47 to 07:58:05 UT for a total of 6918 seconds. All observations were made in the Long Wavelength Spectrometer's Fabry-Perot (LWS/FP) mode (Clegg {\\it et al,} 1996). The instrument's roughly circular field of view was centered on the epoch 2000 coordinates $5h\\;35m\\;14.2s,\\ -5^{\\circ}\\; 22'\\; 23.3\"$. We obtained data on eight H$_2$O lines. These detections required the use of four different detectors, each having its own roughly elliptical beam size --- $70\\times 68$ arc seconds at 125 and 121 $\\mu$m, $77\\times 71$ at 99.5, 95.6 and 90 $\\mu$m, $82\\times 76$ at 83 and 82 $\\mu$m, and $83\\times 79$ at 72 $\\mu$m. The fields of view, therefore, differed by $\\pm 19\\%$. In the wavelength range from 72 to 125 $\\mu$m the LWS/FP's spectral resolving power gradually increases from $\\sim 7,000$ to a peak of $\\sim 9,800$ at 95 $\\mu$m, before slowly dropping to $\\sim 9,500$ at 125$\\mu$m. We repeatedly stepped the Fabry-Perot over a wavelength range that spanned 5 resolution elements on either side of each line. This made for a total of 11 resolution elements, each of which was sampled at 8 equally spaced positions within the element, in what is designated as the ``rapid scanning\" mode of operation, each measurement lasting only 0.5 sec. Strong signals were observed in all eight lines (see Figure 1). ", + "conclusions": "Molecular shocks are important not only for their intrinsic interest, but also because current views assume that star formation may well be triggered by shock compression followed by rapid cooling. The cooling needs to be rapid in order to prevent a quasi-elastic bounce which would permit shock-compressed regions to rebound to their original dimensions. If a shocked region is able to radiate away a substantial fraction of its energy during the traversal time of the shock, it will remain compressed. Even if it is not sufficiently dense at this stage to enter protostellar collapse, it will be poised to contract further if subjected to subsequent shocks. We can estimate the integrated water vapor line flux from Orion summed over all transitions and compare it to the power radiated away by H$_2$ and CO. To estimate the total water vapor emission we can sum the flux both from lines that we have observed and lines whose strengths we infer by applying the KN model. This leads us to deduce a total water vapor emission of $\\sim 0.11\\,$erg cm$^{-2}$ s$^{-1}$ sr$^{-1}$ normalized to a 75 arc second diameter solid angle. The beam-averaged CO flux from the region is $\\sim 0.13$ erg cm$^{-2}$ s$^{-1}$ sr$^{-1}$ obtained from the observational data cited in Stacey {\\it et al.}( 1983), normalized to an assumed beam size of 60 arc seconds --- slightly smaller than the beam size used for our H$_2$O observations. This value could, however, drop by $\\sim 35\\%$ if the shocked region subtended a substantially smaller solid angle than our 75 arc second beam. The H$_2$O and CO cooling of the shocked region are therefore comparable; but both are an order of magnitude lower than the total H$_2$ emission." + }, + "9802/astro-ph9802036_arXiv.txt": { + "abstract": "The current evidence for morphologically peculiar galaxy populations at high-redshifts is outlined. After describing various techniques which can be used to quantify the importance of ``morphological K-corrections'', and to objectively classify galaxy morphology in the presence of these biases, it is concluded that observational biases are not sufficient to explain the increase in the fraction of peculiar galaxies on deep HST images. A new technique is then described which models the spatially resolved internal colors of high redshift galaxies, as a probe of the processes driving galaxy evolution. This ``morphophotometric'' approach investigates directly the evolutionary history of stellar populations, and is a sensitive test of the mechanisms through which galaxies build up and evolve in the field. As a case study, we analyse several ``chain galaxies'' in the Hubble Deep Field. These chain galaxies are shown to be protogalaxies undergoing their first significant episodes of star-formation, and not simply distant edge-on spirals. ", + "introduction": "Recent work from deep imaging \\cite{Griffiths:1994,Glazebrook:1995,Driver:1995,Abraham:1996a,Abraham:1996b,Giavalisco:1996} and spectroscopic \\cite{Lilly:1995,Cowie:1995,Lilly:1996,Ellis:1996} surveys has shown that much of the rapidly evolving faint blue galaxy population \\cite{Broadhurst:1988,Koo:1992,Lilly:1995} is comprised of morphologically peculiar galaxies. These systems may be luminous counterparts to local irregular galaxies, tidally disturbed systems, or perhaps members of entirely new classes of objects with no local counterpart. Another possibility is that these morphologically peculiar systems are simply ``ordinary'' galaxies whose strange appearance is simply a result of their being observed in the rest-frame ultraviolet (a ``morphological K-correction''), where we know little about the appearance of the galaxy population. This distinction between intrinsic and apparent peculiar galaxies lies at the heart of this meeting. In this article several lines of evidence are reviewed which suggest that the bulk of the morphological peculiarities seen in distant galaxies are intrinsic to these systems, and not simply the product of rest-frame bandshifting. In the final section of this article preliminary results from a new line of evidence are presented, focussing on the ``chain galaxy'' population as a case study. ", + "conclusions": "" + }, + "9802/astro-ph9802200_arXiv.txt": { + "abstract": "The variability observed in many complex gamma-ray bursts (GRBs) is inconsistent with causally connected variations in a single, symmetric, relativistic shell interacting with the ambient material (``external shocks''). Rather, either the central site must produce $\\sim 10^{50}$ erg s$^{-1}$ for hundreds of seconds (``internal shocks'') or the local spherical symmetry of the shell must be broken on an angular scale much smaller than $\\Gamma^{-1}$ where $\\Gamma$ is the bulk Lorentz factor for the shell. The observed variability in the external shock models arises from the number of causally connected regions that (randomly) become active. We define the ``surface filling factor'' to be the ratio of the area of causally connected regions that become active to the observable area of the shell. From the observed variability in 52 BATSE bursts, we estimate the surface filling factor to be typically $\\sim 5 \\times 10^{-3}$ although some values are near unity. We find that the surface filling factor, $f$, is $\\sim 0.1 \\Delta T/T$ in both the constant $\\Gamma$ phase (which probably produces the GRB) and the deaccelerating phase (which probably produces the x-ray afterglows). Here, $\\Delta T$ is a typical time scale of variability and $T$ is the time since the initial signal. We analyze the 2 hr flare seen by ASCA 36 hr after the GRB and conclude that the surface filling factor must be small ($10^{-3}$) in the x-ray afterglow phase as well. Compared to the energy required for an isotropic shell, $E_{\\rm iso}$, explanations for low surface filling factor can either require more energy ($f^{-1}E_{\\rm iso} \\sim 10^{56}$ erg) or less energy ($({\\Delta T \\over 4T})^2E_{\\rm iso} \\sim 10^{49}$ erg). Thus, the low filling factor cannot be used as a strong argument that GRBs must be internal shocks. ", + "introduction": "Gamma-ray burst (GRB) spectra often extend to very high energies with no indication of attenuation by photon-photon interactions. This implies substantial relativistic bulk motion of the radiating material with Lorentz factors in the range of $10^2$ to $10^3$. At cosmological distances, GRBs require an energy reservoir on the order of $10^{52}$ erg. The likely sources of such a reservoir would be the rest mass of a compact object released during a merger (e.g., either neutron star -- neutron star or neutron star -- black hole). However, most detailed calculations of mergers occur on time scales of less than 1 s, much less than the observed durations of GRBs (often $10^2$ s and sometimes $10^3$ s). Two classes of models have arisen that explain various (but not all) aspects of the observations. In the ``external'' shock models (\\cite{mr93}), the release of energy during the merger is very quick, and a relativistic shell forms that expands outward for a long period of time ($10^5$ to $10^7$ s). At some point, interactions with the external medium (hence the name) cause the energy of the bulk motion to be converted to gamma-rays. Although the shell might produce gamma-rays for a long period of time, the shell keeps up with the photons such that they arrive at a detector over a relatively short period of time. If the shell has a velocity, $v = \\beta c$, with a corresponding bulk Lorentz factor, $\\Gamma = (1-\\beta^2)^{-1/2}$, then photons emitted over a period $t$ arrive at a detector over a much shorter period, $T = (1-\\beta)t = t/(2\\Gamma^2)$. Although this model is consistent with the short energy release expected for a merger and the observed long time scale of GRBs, we have argued that it cannot explain the long complex time histories except under extreme conditions (\\cite{fmn96}). In particular, we argue from kinematic considerations that the long gaps often seen in GRBs are inconsistent with local spherical symmetry, that the rapid time variability implies that only a small fraction of the shell becomes active, and that the observed average envelope of emission is inconsistent with that expected from a relativistic shell. These arguments are contained in \\cite{fmn96n,fs97n}, and \\cite{frs98n}. In addition, \\cite{dar97n}, \\cite{sp97bn}, and \\cite{kps97n} argue from hydrodynamic considerations that the external shock model cannot provide the energy or timescales observed in GRBs. The alternative theory is that a central site releases energy in the form of a wind or multiple shells over a period of time commensurate with the observed duration of the GRB (\\cite{rm94}). The gamma-rays are produced by the internal interactions within the wind; hence these scenarios are often referred to as internal shock models. These models have two weaknesses: first, there is a concern that internal shocks are rather inefficient (although, see \\cite{kps97}), and second, the long, complex time history of a GRB must be postulated at the central site. On the other hand, the free parameters associated with the multiple shells can probably explain any observed time history (\\cite{kps97}). The need to explain long time histories was further complicated by the discovery of the x-ray afterglows lasting hours (\\cite{cos97}), the optical afterglows lasting weeks to months (\\cite{metzger97}), and the radio afterglows lasting many months (\\cite{frail97}). These extended time scales appear too long to have been produced by a lingering central site made by a merger. In addition, the observed power law decay is expected from many external shock models (\\cite{wrm97,danr97,tavani97,wkf98,pm98,mrw98,rm98,spn98}). \\cite{ps97n} suggested that the initial gamma-ray phase is due to internal shocks from a relativistic wind (or multiple shells) that merge into a single relativistic shell which then produces the afterglows in a manner similar to the external shock models. This model avoids the difficulty of explaining the rapid variability of the gamma-ray phase with a single shell while retaining the long time scale capability of the single shell for the afterglows. The purpose of this paper is to demonstrate that the external shock model cannot utilize the full surface of shells. Fundamentally, it arises because of the relatively rapid time variability of the gamma-ray phase. \\cite{sp97a} use this as an argument that external shock models would require too much energy. However, this should not be considered a strong argument that the GRB phase is not external shocks since there are several scenarios that can utilize only a portion of the shell's surface. We also analyze the time variability recently observed by ASCA in the x-ray afterglow of GB970828 (\\cite{yoshida98}) and show that its time variability implies low surface utilization in the x-ray afterglow phase as well. ", + "conclusions": "The rapid variations in GRB time histories imply emitting entities the size of $\\Delta R_\\perp \\sim c\\Gamma\\Delta T_p$. Assuming a {\\it single} expanding shell, these entities must form on a much larger surface, $\\sim c\\Gamma T$. We have defined $f$ (given by eq. [\\EFFEQ]) to be the fraction of the surface of the shell which becomes active, that is, generates emitting entities. A crude estimate of $f$ is $\\sim 0.1 \\Delta T_p/T$. Because both the classic gamma ray burst phase and the recently discovered x-ray afterglow phase can have flares with $\\Delta T_p /T \\sim 0.05$, both time periods can suffer from an inefficient utilization of a single relativistic shell. In the case of the gamma-ray phase, there is a simple explanation: the gamma-ray source is not a single relativistic shell, but rather, reflects activity at some central site that produces $\\sim 10^{50}$ erg s$^{-1}$ for up to several hundred seconds. Other observations, particularly gaps seen in the gamma-ray time histories (\\cite{frs98}), also argue for a central engine. The x-ray afterglow, however, is widely interpreted as being a single relativistic shell (\\cite{wrm97,danr97,wkf98,spn98,rm98}). Therefore, we will concentrate our discussion on how to accommodate flares during the x-ray afterglow. These arguments could also be applied to the gamma-ray phase. Of the two examples of deviations from a power law in the x-ray afterglow, the \\cite{yoshida98n} event (GB970828) is much more restrictive than the \\cite{piro97n} event (GB970508). GB970828 has a narrow flare whereas GB970508 appears to be roughly a step function. A step function is similar to a FRED-like deviation which is approximately what one would {\\it expect} if most of the observable surface of a relativistic shell interacts with a large ISM cloud (i.e., $\\Delta T/T$ should be $\\sim 0.25$, \\cite{fmn96}). Thus, GB970828 might be an afterglow with a low surface filling factor whereas GB970508 has a high surface filling factor. A common misconception is that one can just use an ISM cloud that covers most of the shell's surface and, therefore, make a flare that substantially changes the emission. This does not work because the curvature of the expanding shell prevents the shell from engaging the cloud instantaneously. Rather, the portion of the shell at $\\theta \\sim \\Gamma^{-1}$ requires a time $R(1-\\cos\\theta)/v$ longer to reach the cloud. Even if the cloud happens to have a concave shape such that the shell reaches the cloud simultaneously over a wide range of angles, the resulting photons at $\\theta \\sim \\Gamma^{-1}$ must travel further to the detector resulting in emission that is delayed by $R(1-\\cos\\theta)/c$. Indeed, the delay of the photons due to the curvature is identical to the scenario of a shell that expands without producing photons for a long time and then emits only over a small range of times (the $P=P_0\\delta(t-t_0)$ case in \\cite{fmn96}). The result is a FRED-like flare (like GB970508) whose duration, $\\Delta T$, is about $T$, rather than a small fraction of $T$ as indicated by the ASCA observations for GB970828. The shell will have well-defined curvature because all points on it are moving at nearly the same speed, that is, the speed of light. In any case, there is no reason to believe that variations in the ambient material would cause the shell to develop into a plane wave oriented towards the observer such that the photons produced by an interaction with an ISM cloud or a shock would arrive as a short flare. {\\it Only the instantaneous interaction between two plain parallel surfaces oriented perpendicularly to our line of sight can produce a short peak from large surfaces.} We see six possible explanations for flares with small $\\Delta T/T$ that produce large changes in the observed intensity. These explanations are explained below and represented in Figure 5. A key issue is the size of the corresponding energy reservoir required in each case. To intercompare these energy requirements, we use that a typical peak corresponds to an isotropic luminosity, $L_0$, equal to $10^{51}$ erg s$^{-1}$. The total energy reservior is required to have $$ E_{\\rm tot} = L_0 {{\\rm d}\\Omega \\over 4\\pi} \\epsilon^{-1} T_{\\rm Dur} \\eqno(\\neweq) $$ where $T_{\\rm Dur}$ is the average duration of the event (say, 30 s) and $\\epsilon$ is the fraction of the total energy that is converted into gamma-rays (say $10^{-1}$). If the emission is isotropic, then the total energy is about $E_{\\rm iso} = 3 \\times 10^{53}$ erg. Since we cannot observe the energy released at angles much greater than $\\sim\\Gamma^{-1}$, any model can reduce the required energy reservior to ${{\\rm d}\\Omega \\over 4\\pi}E_{\\rm iso}$. Although the energy requirement can be made small, a single shell that emits nearly uniformly over \\domega\\ (see Fig. 5a) is not consistent with the observed time structure since a uniform shell always results in a FRED-like shape. The following scenarios can produce bursts with peaks with small $\\Delta T/T$: \\begin{enumerate} \\item{ The simplest way to have short powerful flares is to invoke a central engine. In a central engine, there are multiple releases of energy at the central site over a time scale commensurate with the observed duration of the GRB. Each release makes a peak in the time history. Indeed, small $\\Delta T/T$ and gaps are the key kinematic reasons why the GRB phase is likely to be a central engine (\\cite{fmn96,frs98}). } \\item{ It is possible for a large fraction of the surface of the shell to be active, and the flare region might still be small, because $\\epsilon$ or the energy content of the shell varies strongly as a function of position on the shell. That is, the photons cm$^{-2}$ s$^{-1}$ at the flare site must be larger than the rest of the shell by a factor of $\\sim f^{-1}$ (see Fig. 5b). For example, some photon production rates scale as $\\Gamma^3$, so a single region with a $\\Gamma$ that is 10 times larger than other regions could produce emission that doubles the overall count rate. This would give the appearance that $f$ is $\\sim 10^{-3}$. If the energy content of the shell is constant as a function of angle and $\\epsilon$ varies, then the required total energy is about $f^{-1}{{\\rm d}\\Omega \\over 4\\pi}E_{\\rm iso}$. If $\\epsilon$ is constant, and the energy content of the shell varies, then the total required energy is ${{\\rm d}\\Omega \\over 4\\pi}E_{\\rm iso}$. } \\item{ The third explanation is that the local spherical symmetry of the shell is broken into many emitting regions on a scale much smaller than $\\Gamma^{-1}$ where only a fraction $f$ of the surface converts its energy into photons. The regions are small enough ($\\Gamma c \\Delta T$) to make peaks with small $\\Delta T/T$ (see Fig. 5c). Here it is assumed that all regions of the shell has equal energy content but much of it is wasted because most regions do not produce gamma-rays. Thus, this scenario requires an energy reservoir of $f^{-1}{{\\rm d}\\Omega \\over 4\\pi}E_{\\rm iso}$, that is, $f^{-1} \\sim 10^3$ {\\it more} energy than a simple shell. If \\domega\\ is large, this scenario could require $E_{\\rm tot}$ to be as large as $10^{56}$ erg. \\cite{sp97an} invoke this explanation and the resulting need for more energy to argue against a single relativistic shell. That arguement should not be considered particularly strong since we have several scenarios that do not waste the energy and/or \\domega\\ could be small. } \\item{ The fourth explanation is that the local spherical symmetry is broken because the outflow of material is in narrow fingers that occupy only a fraction $f$ of the shell's surface (see Fig. 5d). This is basically the same geometry as the previous case except those regions that do not emit gamma-rays do not have energy content. Since no energy is wasted in non-emitting regions, this scenario requires an energy reservoir of only ${{\\rm d}\\Omega \\over 4\\pi}E_{\\rm iso}$, that is, about what one would expect from a simple shell. } \\item { The fifth explanation is that \\domega~is very small, so small that the emission lasts only $\\Delta T_p$ even if the entire \\domega~region is gamma-ray active. This requires \\domega~to be $\\sim ({\\Delta T_p \\over 4T})^2$. The requried total energy is correspondingly small: $({\\Delta T_p \\over 4T})^2E_{\\rm iso} = 10^{49}$ erg. This explanation has one narrow ejecta that needs to go through a gamma-ray emitting phase for each peak in the GRB time history. In the above explanation (number 4), there are many narrow ejecta, so each one only has to become gamma-ray active only a few times to make a GRB time history with hundreds of peaks. } \\item{The last explanation attributes the duration of the GRB to the spreading out in time of the emission because of different conditions as a function of $\\theta$ (see Fig. 5f). This is what we called a ``thick shell with substructure'' in \\cite{fmn96}. It would have been better to call it a ``thick emitting region with substructure'' since the actual emitting shell must always be thin (cf. \\cite{fmn96}). Effectively, $f$ is unity but it appears to be small because the assumptions involved in equation (\\AREASHELL) are not valid. Let $t_i$ be the time that the $i$-th section of a shell at $\\theta_i$ expands before becoming active. Let $\\beta_i$ be the speed of the material at $\\theta_i$. The relative time that the emission from two different regions ($i$ and $j$) arrive at our detector is $$ T = T_i - T_j = t_i(1-\\beta_i\\cos\\theta_i) - t_j(1-\\beta_j\\cos\\theta_j)~~. \\eqno(\\neweq) $$ Explanations 3 and 4, above, assume that all $t_i$'s and $\\beta_i$'s are equal such that we can estimate the size of the shell from it's FRED-like appearance and peaks arrive at different times because they come from different $\\theta_i$'s. Of course, most bursts are not FRED-like. In this explanation, all $\\theta_i$ regions become active (so $f$ =1), but the peaks that would have added up to make the FRED-like envelope are scattered in time because of the different $t_i$'s and $\\beta_i$'s. To disrupt the FRED-like shape, $T_i$'s must vary by a factor of at least 2 from what they would be if everything was constant with $\\theta$. Consider the effect of different times of emission (i.e., different $t_i$'s but similar $\\beta_i$'s). The $t_i$'s would have to vary by about a factor of 2 to change $T_i$'s by a factor of 2. Consider constant $t_i$'s but different $\\beta_i$'s. To spread out out the emitting regions so their distance from us varies by about $T$ requires that $(\\beta_j-\\beta_i)ct=T$ or that the radius of the shell is $$ ct = {2\\Gamma_i^2\\Gamma_j^2 \\over \\Gamma_j^2-\\Gamma_i^2}cT \\sim 2\\Gamma_i^2cT \\eqno(\\neweq) $$ The $\\Gamma_i$ and $\\Gamma_j$ would have to vary by about $\\sqrt{2}$. Such differences could arise from different baryon loading as a function of $\\theta$. There are two observations that argue against this explanation. First, even the bursts with FRED-like envelopes show low filling factors (see the solid squares in Fig. 3). Second, if the $\\Gamma_i$'s differ by $\\sqrt{2}$, then peak near the end of the burst should have widths that are at least a factor of 2 wider. Bursts normally do not appear to have such a trend. This explanation requires an energy reservoir like a uniform shell: ${{\\rm d}\\Omega \\over 4\\pi}E_{\\rm iso}$. } \\end{enumerate} In conclusion, both the gamma-ray phase and the x-ray afterglow show similar rapid variability, both have values of $f \\sim 0.1\\Delta T/T$ that are similar. Such variability implies that either GRBs are central engines, or that the local spherical symmetry is broken on a scale much smaller than $\\Gamma^{-1}$. In the last case, the structure can either be due to variations in $\\Gamma$ as a function of $\\theta$ (perhaps due to small scale variations in the baryon loading, explanation 6 above), or due to only a fraction $f$ of the shell becoming active (explanations 3, 4, 5 above). The required total energy is $\\sim f^{-1}E_{\\rm iso}$, $E_{\\rm iso},$ and $({\\Delta T_p \\over 4T})^2E_{\\rm iso}$ for explanations 3, 4, and 5, respectively. Using typical values, the total energy can vary from $10^{56}$ to $10^{49}$ erg. We note that explanation 5 (a jet much narrower than $\\Gamma^{-1}$) is the only explanation that can also easily explain gaps in the time history." + }, + "9802/astro-ph9802115_arXiv.txt": { + "abstract": "We compute the evolution of the space-dependent mass distribution of galaxies in clusters due to binary aggregations by solving a space-dependent Smoluchowski equation. From the solutions we derive the distribution of intergalactic distance for different ranges of mass (and of corresponding magnitude). We compare the results with the observed distributions, and find that the different degrees of luminosity segregation observed in clusters are well accounted for by our merging model. In addition, the presence of luminosity segregation is related to dynamical effects which also show up in different but connected observables, such galaxy velocity profiles decreasing toward the center and X-ray measured $\\beta$-parameters smaller than 1. We predict both luminosity segregation and the observables above (being a product of binary aggregations) to be inversely correlated with the core radius and with the galaxy velocity dispersion; we discuss how the whole set of predictions compares with up-to-date observations. ", + "introduction": "The dynamical evolution of galaxy clusters is currently believed to go through two major phases: in the first, usually referred to as violent relaxation (Lynden-Bell 1967), the evolution is controlled by a collective potential and results in a Maxwell velocity distribution of galaxies; in the second, the dynamics is dominated by two-body processes, and binary (both elastic and inelastic) collisions drive the evolution. In fact, in this latter phase, for ordinary galaxy sizes and separations the collision time scale is much less than the Hubble time. Although a complete theoretical description of the two-body phase of dynamical evolution of clusters is still lacking, observations, N-body simulations and computation based on statistical methods (Monte Carlo and Fokker-Plank simulations) have concurred in enlightening many dynamical properties of clusters in this stage. Such properties are characterized by a large cluster-to-cluster variance, and include the following: the presence of a {\\it velocity bias} $b_v^2=\\langle v^2 \\rangle/\\sigma^2 <1$ of the galaxy velocity dispersion $\\langle v^2\\rangle^{1/2} $ with respect to the dark matter's $\\sigma$ (see N-body simulations by Carlberg \\& Dubinski 1991; Evrard, Summers \\& David 1994; Katz \\& White 1993; Carlberg 1994; Summers, Davis $\\&$ Evrard 1995); galaxy {\\it velocity dispersion profiles decreasing toward the cluster center} (see observations by Kent \\& Sargent 1983; Sharples, Ellis, \\& Gray 1988; Girardi et al. 1996); {\\it mass segregation}, i.e., the tendency of more massive galaxies to be located near the cluster center (see simulations by Roos \\& Aarseth 1982; Farouki, Hoffman \\& Spencer 1983) with the associated luminosity segregation (observed in several clusters, see Rood et al. 1972; Oemler 1974; White 1977; Dressler 1978; Quintana 1979; Sarazin 1980; Kent \\& Gunn 1982; Oegerle, Hoessel \\& Ernst 1986; Binggeli, Tammann \\& Sandage 1987; Dominguez Tenreiro \\& Del Pozo-Sanz 1992; Stein 1996). Due to the large variance observed in the above effects, appreciable uncertainties exist about their dependence on the characteristics of the system, although a general trend in the sense of larger effects for smaller clusters might be inferred from the data, when only relaxed cluster (with no prominent substructure or asphericity) are considered. A complete description of the post-virialization phase of galaxy cluster should be able to connect all the above effects and to explain the observed variance in terms of dynamical properties of clusters. In previous papers (Fusco-Femiano $\\&$ Menci 1995, hereafter Paper I; Menci \\& Fusco-Femiano 1996, hereafter Paper II) we showed that the loss of kinetic energy in inelastic galaxy collisions (binary aggregations) in clusters with $\\sigma\\lessim 900$ km/s can significantly change the velocity distribution of galaxies. The model not only simultaneously accounts for the velocity bias and for centrally decreasing velocity dispersion profiles, but also predicts a correlation of such effects with the shape (the core radius) and the depth (the dark matter velocity dispersion) of the cluster potential wells in good agreement with observations. In addition, the aggregation model can succesfully connect the above dynamical effects to other observed properties of galaxy clusters (Paper I), such as the $\\beta$-parameter (expressing the ratio of the galaxy orbital specific energy to the specific energy of the X-ray emitting plasma) and the Butcher-Oemler effect (see Cavaliere \\& Menci 1993). Thus, aggregations seem to constitute a leading mechanism in the post-virialization phase of clusters with velocity dispersion $\\lessim 900$ km/s, while in larger clusters, they are highly suppressed due to the large galaxy relative velocities (see numerical results by Richstone \\& Malmuth 1993). To further assess the role of galaxy merging in the two-body dynamical phase, we address here the problem of luminosity segregation (hereafter LS). This is expected to be generated when aggregations are effective, since merging builds up larger galaxies mainly in the central regions, where the larger density favours binary aggregation. Thus, we extend our treatement of galaxy inelastic collisions, based on the solution of a collisional Boltzmann-Liuoville equation, to include position-dependent mass spectra of interacting galaxies. The prediction of our model will be compared with observational results, focussing on the correlation of the segregation effects with the properties of the clusters such as the richness, the density distribution and the velocity dispersion. Finally, we shall show how the merging model connects LS with the dynamical effects discussed above and with their X-ray counterparts. The paper is organized as follows. In sect. 2 we discuss the collisional Boltzmann equation and describe our solutions for the evolution of the position-dependent mass distribution. In Sect. 3a we decribe the standard method we use for the comparison with data, based on the auto-correlation function for galaxies of different mass. The comparison is performed in Sect. 3b. Sect. 4 is devoted to discussion and conclusion. \\section { Radius-dependent mass distribution from binary aggregations} \\subsection{The Boltzman Equation for Merging Galaxies} The evolution with time $t$ of the distribution $f_t(M,{\\bf r},{\\bf v})$ of interacting galaxies with velocity ${\\bf v}$ and mass $M$ at the position ${\\bf r}$ inside the gravitational potential $\\psi$ of a cluster can be described by the collisional Boltzmann equation. Assuming spherical symmetry, the latter can be written in spherical coordinates ${\\bf r}=(r,\\theta,\\phi)$ and ${\\bf v}=(v_r,v_{\\theta},v_{\\phi})$ for the distribution $f(M,r,v_r,v_t)$ as follows (see, e.g., Saslaw 1985) \\begin{eqnarray} \\lefteqn{ \\partial_t f_t(M,r,v_r,v_t) +v_r\\,{\\partial f_t(M,r,v_r,v_t) \\over\\partial r}+ \\Big({v_t^2\\over r}-{\\partial\\psi \\over\\partial r}\\Big)- 2\\,{v_r\\,v_t^2\\over r}\\,{\\partial f_t(M,r,v_r,v_t)\\over \\partial v_t^2} = } \\nonumber \\\\ & & {1\\over 2}\\, \\int_o^M\\,dM'\\int_{-\\infty}^{\\infty}dv_r'\\,\\int_0^{\\infty}\\,dv_t'^2 f_t(M',r,v_r',v_t')\\,f_t(M-M',r,v_r'',v_t'')\\,\\Sigma(M',M-M',v_{rel}) \\,v_{rel}+ \\nonumber \\\\ & & -\\int_0^{\\infty}dM'\\,\\int_{-\\infty}^{\\infty}dv_r'\\,\\int_0^{\\infty}\\,dv_t'^2\\, f_t(M,r,v_r,v_t)\\,f_t(M',r,v_r',v_t')\\,\\Sigma(M',M',v_{rel})\\,v_{rel} \\end{eqnarray} where $v_t^2\\equiv v_{\\theta}^2+v_{\\phi}^2$ is the square tangential component of velocity, and the gravitational cross section for interactions $\\Sigma$ depends on the relative velocity ${\\bf v}_{rel}\\equiv {\\bf v}'-{\\bf v}''$. The velocity ${\\bf v}''$ in the first integral in eq. (1) is related to ${\\bf v}$ and to ${\\bf v}'$ by the requirement of momentum conservation $M'\\,{\\bf v}'+(M-M')\\,{\\bf v}''=M\\,{\\bf v}$. Here we have assumed that the galaxies do not gain or loose mass via processes other than merging. To obtain a fully self-consistent description, eq. (1) should be complemented with the Poisson equation for the gravitational potential $\\psi$ with a source term $\\int dM\\,dv_r\\,dv^2_t\\, f_t(M,r,v_r,v_t)$, which includes the galaxy distribution itself. However, since we are interested in the post-virialization phase of cluster evolution where the potential is essentially fixed, we shall assume a King potential $\\psi (r)$ and follow the evolution of the galaxy mass distribution at different radii. Since our aim is to probe the effectiveness of interaction in producing mass segregation, we indroduce some approximations (a discussion on them is given in the final section). First, we assume the velocity distribution to be independent on the mass and on the spatial distribution of galaxies, so that the distribution in eq. (1) can be factorized into a velocity distribution $p({\\bf v})$ and a position-dependent mass distribution $N(M,r,t)$. Such approximation does not actually hold (see Paper I) but, as we discuss in the final Section, for our purpose in the present paper it is a {\\it conservative} assumption. Second, we assume that the radial and tangential velocity distribution are mutually independent and both normally distributed. In this case, integration of eq. (1) over velocities leads to the following position-dependent Smoluchowski equation \\begin{eqnarray} \\partial_t N(M,r,t) & = & {1\\over 2}\\int dM'\\,N(M',r,t)\\,N(M-M',r,t) \\langle\\Sigma (M',M-M')\\,v_{rel}\\rangle + \\nonumber \\\\ & & -\\int dM'\\,N(M,r,t)\\,N(M',r,t)\\langle\\Sigma (M,M')\\,v_{rel}\\rangle \\end{eqnarray} where the average $\\langle\\rangle$ is over the velocity distribution. The cross section is given by (Saslaw 1985): $\\Sigma(M,M')=\\epsilon(v_{rel}/v_g)\\, \\pi\\,(r^2+r'^2)\\, \\Big[1+{v_g^2\\over \\,v_{rel}^2}\\Big]$ where $r$ and $r'$ are the radii of the interacting galaxies (proportional to $M^{2/3}$) and $v_g\\propto G(M+M')/R$ is the escape velocity at closest approach $R\\approx (r+r')$. The efficiency $\\epsilon$ is determined from N-body results (see Richstone \\& Malmuth 1993) and is zero when $v_{rel}\\gtrsim 3 v_g$, so that aggregations are highly suppressed in very rich clusters. It is convenient to express all quantities in terms of the {\\it adimensional} mass $m\\equiv M/M_*$ normalized to the characteristic mass $M_*$ (that corresponds to a galaxy with characteristic luminosity $L_*$). From $r\\sim(M/\\rho)^{1/3}$, the relation $r^2=r_{g*}^2\\,m^{2/3}$ follows. Then, the cross section reads \\begin{equation} \\Sigma(m,m')=\\epsilon(v_{rel}/v_g)\\,\\pi\\,r_{g*}^2 \\,(m^{2/3}+m'^{2/3})\\, \\Big[1+(m^{2/3}+m'^{2/3})\\,v^2_{g*}/v^2_{rel}\\Big] \\end{equation} where $r_{g*}$ and $v_{g*}$ are the radius and the 3-D internal velocity dispersion of a $L_*$ galaxy, respectively. \\subsection{ Initial Conditions} We assume the galaxy distribution to be {\\it initially} (after the cluster formation and virialization) factorized in a mass distribution $P(m)$ times a King spatial profile, which will be subsequently mixed up by the 2-body dynamical evolution. Then \\begin{equation} N(M,r)_{t=0}={n_o\\over (1+x^2)^{3/2}}\\,P(m)~, \\end{equation} where we take for $P(m)$ the Press \\& Shechter shape $P(m)=m^{a-2}\\,e^{-b^2\\,\\delta_c^2\\,m^{2a}/2}$ (the index $a$ depends on the spectrum of cosmological perturbations and is in the range 0 - 0.3 at the scale of galaxy clusters) and $x=r/r_c$ is the distance from the cluster center in units of the core radius $r_c$ of the King profile. The constant $n_o$ is taken as to yield the total number $N_{tot}$ of galaxies inside the cluster virial radius $R_v$ (from the virial theorem $R_v=G\\,M/3\\sigma$); thus \\begin{equation} n_o={N_{tot}\\over 4\\,\\pi\\,r_c^3\\,I_R\\,I_M}~, \\end{equation} where $I_R\\equiv \\int_0^{R_v/r_c}dx\\,x^2/(1+x^2)^{3/2}$ and $I_M\\equiv \\int_0^{\\infty}\\,dm\\,P(m)$ are the adimensional integrals of the initial spatial and mass distributions, respectively. \\subsection{Numerical Solutions} To integrate eq. (2) we first write it in a completely adimensional form for the normalized $r$-dependent mass distribution $n_t(m,r)\\equiv N(m,r,t)/n_o$. We define the adimensional time variable $\\tau\\equiv t/t_{cr}\\approx 2\\,10^9\\,{\\rm yr}\\, \\big[R_v/1\\,{\\rm Mpc}\\big]\\,\\big[\\sigma/10^3\\,{\\rm km/s}\\big]^{-1}$ in terms of the the cluster crossing time $t_{cr}\\equiv 2\\,R_v/\\sigma$. The adimensional velocities $\\tilde{v}\\equiv v/\\sigma$ are normalized to the dark matter velocity dispersion $\\sigma$. The corresponding adimensional interaction rate $\\eta (m,m')=n_o\\,\\Sigma\\,R_v\\,\\tilde v$ can be computed from eq. (3) and (4). Then, the Smoluchowski eq. for the normalized mass distribution $n_{\\tau}(m,r)\\equiv N(M,r,\\tau)/n_o$ can be recast in the form \\begin{mathletters} \\begin{eqnarray} \\partial_{\\tau} n_{\\tau}(m,r) & = & {1\\over 2}\\,\\int_0^m\\,dm'\\,n_{\\tau}(m',r)\\,n_{\\tau}(m-m',r)\\, \\langle\\eta(m',m-m')\\rangle+ \\nonumber \\\\ & & -\\int_o^m\\,dm'n_{\\tau}(m,r)\\,n_{\\tau}(m',r)\\,\\langle\\eta(m,m')\\rangle \\\\ \\eta(m,m') &= &{1\\over 2}\\, {R_v\\over r_c}\\,{r_{g*}^2\\over r_c^2} \\,{N_{tot}\\over I_R\\,I_M}\\,\\tilde v_{rel} \\,(m^{2/3}+m'^{2/3})\\, \\Big[1+(m^{2/3}+m'^{2/3})\\,\\tilde v^2_{g*}/\\tilde v^2_{rel}\\Big] . \\end{eqnarray} \\end{mathletters} From the form (6) it is evident how (for constant $R_v/r_c$ ratio) the effect of aggregation is larger for clusters with small core radius (galaxies in the center are more concentrated) and with a larger number of galaxies $N_{tot}$. The average of the aggregation rate in eq. (6b) \\begin{equation} \\langle \\eta\\rangle\\equiv \\int\\,d\\alpha \\int_o^{ {|\\bf\\tilde{v}_1}-{\\bf\\tilde{v}_2}| =3\\,\\tilde{v}_g} d\\tilde{v}_1\\,\\tilde{v}_1^2\\,p(\\tilde{v}_1)\\,d\\tilde{v}_2\\, \\tilde{v}_2^2\\,p(\\tilde{v}_2)\\, \\eta({\\bf\\tilde{v}_1}-{\\bf\\tilde{v}_2}|/\\tilde{v}_g)~, \\end{equation} is over the velocities $\\tilde{v}_1$ and $\\tilde{v}_2$ (normalized to the dark matter velocity dispersion $\\sigma$) of galaxies colliding with relative angle $\\alpha$; the condition $|{\\bf\\tilde{v}_1}-{\\bf\\tilde{v}_2}| \\leq 3\\,\\tilde{v_g}$ accounts for the efficiency $\\epsilon (v_{rel}/v_g)$. We assume the distribution of velocities $p(\\tilde{v})=(1/2\\pi)^{-3/2}\\,e^{-\\tilde{v}^2/2}$ to be Gaussian, as expected after violent relaxation (Lynden-Bell 1967). Note that, for clusters with $\\sigma\\lesssim 900$ km/s, eq. (7) yields significant averaged aggregation rates $\\langle \\eta\\rangle$ assuming a 3-D internal velocity dispersion $v_{g*}= 300$ km/s for an $L_*$ galaxy with $r_{g*}=60\\,h^{-1}$ kpc \\footnote{In the text we adopt $h=0.5$ for the Hubble constant $H_o=100\\,h$ km/s/Mpc. }. The adopted value of $v_{g*}$ correspond to a circular velocity of $\\approx 220$ km/s; such value is consistent with that derived from the Faber-Jackson relation for an $L_*$ galaxy, and with the measurements by Tonry \\& Davis (1981); Dressler (1984); Dressler (1987). The adopted value of $r_{g*}$ (which refers to the dark halo of an $L_*$ galaxy) is a conservative one, when compared with observational results from absorption lines, measured by Steidel (1995); Lanzetta et al. (1995); Barcons, Lanzetta \\& Webb (1995). Equation (6) is integrated up to $\\tau=5$ with time increments $\\Delta \\tau=1/500$ and mass step $\\Delta m=1/500$, from a minimum mass $m_{inf}=10^{-2}$ to a maximum mass $m_{sup}=10^2$ (integrating up to larger times do not affect sensitively our results). When aggregations are effective, the final mass distribution will be changed from the initial one only in the central core where the galaxy density is larger and binary aggregations are favoured. Thus, in the core larger galaxies will form via binary merging, while the initial mass distribution remains unchanged in the outer regions. The evolution of the mass distribution at different radii in a typical cluster (with $N_{tot}=1000$, $\\sigma=800$ km/s and $r_c=250\\,h^{-1}$ kpc) is shown in figure 1. The distribution flattens in the central region due to the disappearence of small galaxies which aggregate to form larger ones. Since aggregations between galaxies cause a loss of orbital kinetic energy, we expect such effect to be correlated with galaxy velocity dispersions smaller in the central regions, i.e., with velocity profiles falling toward the center (as we discussed in Paper II); this is actually the case, as is shown in figure 2. A further effect is that brighter galaxies (which form from mainly in the central, denser regions) will have smaller relative separations. This latter effect is that observed in several clusters, as we discuss in the next session. The strenght of the above effects depends on the cluster properties, which, in our model, enter only through $N_{tot}$, $\\sigma$ and $r_c$ as is shown by eqs. (6). E.g., for given $N_{tot}$ and $\\sigma$, in clusters with large $r_c$ merging will be less effective (see the merging rate in eq. 6b) because the total number of galaxies is spread out in a larger region. ", + "conclusions": "We have shown that a detailed model for the dynamics of galaxies aggregating in the potential wells of clusters predicts luminosity segregation (LS) effects of the kind observed in real clusters. The correlation of the strenght of the effect with the properties of the clusters predicted in our model is in agreement (see figure 3 and Table 2) with that observed for the limited sample of clusters for which LS has been subject to accurate quantitative measurements. In particular, we predict the effectiveness of aggregations, and hence the degree of LS, to be directly correlated with the number of galaxies in the cluster and inversely correlated with the core radius and with the velocity dispersion (see eq. 6b). The results do not depend on the detail of the initial mass distribution of galaxies in clusters, which we assume to have a Press \\& Shechter form with spectral parameter $a=-2$; such independency is due to the properties of the asymptotic solution of the Smoluchowski equation (describing the evolution of the position-dependent galaxy mass function in our model) and can be traced back to the non-linear nature of such equation. Our results are robust with respect to uncertainties in the input quantities and to the adopted $L(M)$ , as shown in Sect. 3.3. As for our assumption of fixed galaxy velocity distribution, this does not hold when aggregations are effective (see Paper I and II). However, the merging-induced shift of $\\sim (10-15)\\%$ of the velocity dispersion toward smaller values {\\it increases} the efficiency of aggregations, so that our assumption is actually conservative. We have re-run our computation for shifted velocity distributions and found results almost indistinguishable from those presented here. Finally, we stress that no attempt of parameter optimization has been performed. An even better agreement could be found if the cluster parameters were suitable tuned. As for the big picture of the evolution of cluster in the two-body dynamical phase, our model focus on the effects of inelastic collisions not considered in previous works on this subject. In particular, the Fokker-Plank approach by Yepes \\& Dominguez-Tenreiro (1992) considers only elastic collisions by means of a ``mean field'' approximation with the input parameters chosen from a grid of models to show that, within the set of models, it is possible to match the observed segregation effects. Here we solve the collisional Boltzmann equation including inelastic collisions, using, for the input model parameters, the {\\it measured} values. Though the latter are subject to errors, we showed (in Sect. 3.3) that the model is robust to uncertainties $< 12 \\%$ in the parameters. Our results show that inelatic collisions produce appreciable dynamical effects for clusters with one-dimensional velocity dispersion $\\lessim 900$ km/s. Such effects show up in different but connected observables: the velocity bias (due to the {\\it average} loss of kinetic energy in inelastic collisions) $b_v\\approx 0.8-0.9$ can be observed in X-rays in the form of $\\beta$-parameter (see Cavaliere \\& Fusco-Femiano 1976) $\\beta=b_v^2<1$; centrally rising velocity dispersion profiles (due to the {\\it differential} loss of kinetic energy at different radii) are now being measured with great accuracy in the optical (Girardi et al. 1996); different average separations of massive galaxies with respect to the faint ones $b_{\\lambda}\\approx 0.8-0.9$ (see eq. 5), i.e., luminosity segregation (due to the differential mass growth from aggregations at different radii) have been measured in different clusters (see references cited in this paper). We note that, when interpreted in terms of merging-driven evolution, {\\it all} the above effects are predicted to have the {\\it same} dependence on the cluster parameters, i.e., to be larger for clusters with smaller core radii $r_c$ and galaxy velocity dispersions $\\sigma$, although the strength of the $\\sigma$-dependence is mild for LS effects. The observational tests for such predictions are critically affected by the presence of clusters with anisotropies and/or substructures in the observational sample. An inverse correlations of the $\\beta<1$-effect with $r_c$ has been found by, e.g., Jones \\& Forman 1984, while the anti-correlation with $\\sigma$ has been pointed out by Kriss et al. (1983); Jones \\& Forman (1984); Edge \\& Stewart (1991); Bird, Mushotzky \\& Metzler (1995); Jones et al. (1997) but has not been confirmed by the analysis by Lubin \\& Bahcall (1993) and by Girardi et al. (1996). For velocity dispersion profiles decreasing toward the center, the observational situation is still unclear. An anti-correlation with $\\sigma$ has been inferred (see Paper II) from the analysis by Girardi et al. (1996) of a sample of 37 clusters, when clusters with prominent substructures are excluded; however, the detection of such correlation from the data (see also den Hartog, \\& Katgert, 1996) is made difficult by the presence of anisotropies and/or substructures, which can hurt or destroy the effect of the inelastic collision in the two-body relaxation phase. As for the LS, the mild (inverse) dependence on $\\sigma$ of the LS effect from merging makes difficult to observe such correlation. However, in our model, the strong inverse correlation of LS with $r_c$ predicted by our model is confirmed the data analysis by Yepes, Dominguez-Tenreiro \\& Del Pozo-Sanz (1991) on the very limited sample of clusters. LS data for a larger sample of cluster with measured $r_c$ would definitely clarify the issue. Finally, we note that the correlations between different but connected observables predicted by the aggregation model makes it testable already at the present stage of observational capabilities. Further observational progress (in particular in measuring in detail velocity dispersion profiles and X-ray temperatures) can definitely probe the predictions of the merging picture, thus assessing the role of aggregations in the dynamical evolution of clusters." + }, + "9802/astro-ph9802053_arXiv.txt": { + "abstract": "We report on the \\rossi~Proportional Counter Array (PCA) observation of the X-ray burster \\onee~located in the globular cluster Terzan 2. The observation lasted for about 100 kiloseconds and spanned from November 4th to 8th, 1996. The PCA source count rate in the 2-20 keV range was about 470 cts/s. No large spectral variations were observed within our observation as inferred from color-color and hardness-intensity diagram analysis. The persistent X-ray emission shows a high level of noise variability, the so-called High Frequency Noise (HFN) with a fractional Root Mean Squared (RMS) of $\\sim$ 25\\% in the $2 \\times 10^{-3} - 40$ Hz range. The strong HFN together with the hardness of its X-ray spectrum suggest that \\onee~is an ``Atoll'' source which was in its ``Island'' state during the observation. The Fourier Power Density Spectra (PDS) can be modeled in terms of the sum of two ``shot noise'' components for which the shots have a single-side exponential shape (i.e. instantaneous rise and exponential decay). The characteristic shot decay timescales inferred from the best fitting of the PDS are $\\sim 680$ and 16 msec respectively. The two components contribute similarly to the total RMS ($\\sim$ 15\\%). The PDS contains also a third component: a broad and asymmetric peaked noise feature centered at 0.8 Hz. This Quasi Periodic Oscillation-like (QPO) feature contributes at the level of $\\sim$ 10\\% to the total RMS. Neither the shot timescales nor the QPO frequency vary with energy. On the other hand, the integrated RMS of all three components shows a positive correlation with energy up to at least 40 keV. In addition, in the 2--20 keV energy band where the signal to noise ratio is the highest, we found evidence for a high frequency component which shows up in the PDS above 100 Hz. In terms of the shot noise model, a Lorentzian fit of this last component implies a shot decay timescale of $\\sim 1.4$ ms. We also show that \\onee~ has striking timing similarities with the black hole candidate GRO J0422+32 (Nova Persei 1992). Our observation demonstrates that a low frequency QPO simultaneously with a high level of RMS is not a timing signature unique to black holes. This extends the growing list of similarities between Atoll sources and black hole systems. ", + "introduction": "An X-ray burst from a region including the globular cluster Terzan 2 was first observed by OSO-8 \\cite{swank77apjl}. A weak persistent X-ray source was also found by OSO-8, and Grindlay \\cite*{grindlay78apjl} showed that both the position of this source and a revised position of the {\\it Uhuru} source 4U1722-30 \\cite{forman78apjs} were consistent with that of the cluster. Later Grindlay et al. \\cite*{grind80apjl} using the EINSTEIN HRI instrument positioned the X-ray burst source inside the core of Terzan 2. Its name then became 1E1724--3045. Basically reliable spectral observations (EXOSAT and TTM) show that the source has a rather hard power law spectrum (photon index around 2) in the 1-20 keV range with an X-ray flux ranging from 3 to $9 \\times 10^{-10}$ \\ergscm. There is also an indication from EINSTEIN data that at higher flux level ($\\sim 2.1 \\times 10^{-9}$ \\ergscm, i.e. in its high state), the X-ray spectrum softens and is better described by a Bremsstrahlung model with a temperature of $\\sim 6$ keV \\cite{tz2:barret98aa}. Estimates of the cluster distance range from 5.2 to 7.7 kpc \\cite{ortolani97aa}. These values are consistent with the one derived from a type I X-ray burst that showed photospheric expansion \\cite{grind80apjl,tanaka81}. Adopting 7 kpc, the 1-20 keV low state luminosity of the source lies around $\\sim 4 \\times 10^{36}$ \\ergs~indicating that \\onee~ belongs to the class of low luminosity systems. As for the timing properties of the persistent X-ray emission of the source, very little is known. However, during an EXOSAT observation performed with the medium energy (ME) experiment, a QPO at 0.092 $\\pm 0.001$ Hz was found (pulsed fraction $4.8 \\pm 1.2$ \\%, Full Width Half Maximum of $0.053 \\pm 0.002$ Hz). In addition low-frequency noise and the second harmonic were observed \\cite{belli86iau}. Note however that the presence of the QPOs reported in the {\\it IAU} circular was not confirmed in a proceedings paper by the same authors \\cite{belli85esa}. Besides its X-ray properties, \\onee~is remarkable by the fact that it is a source of persistent, though variable, hard X-ray emission ($\\sim 35-150$ keV) as observed with SIGMA \\cite{tz2:barret91apjl,gold93:2ndgro}. \\onee~is in fact the first, and one of the only X-ray bursters (i.e. weakly magnetized neutron star) showing such a persistent emission with a spectrum extending up to $\\sim 100$ keV \\cite{vargas96int,tavani974thcgro}. \\begin{figure}[t] \\vspace{0cm} \\vspace{0cm} \\hspace{0cm}\\centerline{\\psfig{figure=fig1a.ps,width=9.0cm,height=12cm}} \\vspace{0cm} \\vspace{-4cm} \\hspace{0cm}\\centerline{\\psfig{figure=fig1b.ps,width=9.0cm,height=12cm}} \\vspace{-4cm} \\caption{PCA Light curve of 1E1724-3045 in the energy bands 2-5, 5-10, 10-50 keV. The bin time is 512 seconds.} \\label{pcalc} \\end{figure} We proposed \\onee~as a target for a Rossi X-ray Timing Explorer observation for two main reasons: First we wanted to characterize the rapid variability of the source in order to check whether, as expected, \\onee~is an Atoll source. Second we wished to detail the continuum energy spectrum simultaneously from X-rays to hard X-rays. The main focus of this paper is the study of the rapid variability of the source in the frequency range $2 \\times 10^{-3}$ to a few tens of Hz. The search for high frequency QPOs (above $\\ga 300$ Hz) will appear elsewhere. Similarly the spectral analysis of the combined PCA and HEXTE data is the subject of a forthcoming paper \\cite{tz2:barret98aa}. The present paper is organized as follows: In section 2, we present the correlated temporal and spectral analysis using light curves, color-color and hardness-intensity diagrams. In section 3, we present the results of the rapid variability study, including the detailed modeling of the PDS. In section 4, we discuss on the possible origins of the fast timing variability of the source, emphasize the similarities between black hole and neutron star systems, and address the issue about the nature of the low frequency QPO observed. ", + "conclusions": "Our RXTE/PCA observation enabled us to classify \\onee~as an ``Atoll'' source. It was observed in its ``Island'' state for which its X-ray energy spectrum is hard and the PDS displays high frequency noise with a high level of RMS ($\\sim$ 30\\%). The shape of the PDS is complex but can be decomposed in the framework of the shot noise model. In addition the PDS contains a QPO-like feature peaking at $\\sim 0.8$ Hz. Similar PDS properties have already been observed from several black hole candidates, and in particular from \\groj~for which the similarities are striking. A complete timing analysis of the PCA data is in progress with the aim of studying the shot profiles, the origin of the QPO, the spectral evolution along the shots, time lags as a function of energy, etc... \\cite{tz2:olive98aa}. This illustrates well the wealth of the PCA data. Some very valuable information will also be provided by the HEXTE experiment which will for the first time enable us to study the hard power law tail and the aperiodic variability simultaneously from X-rays to hard X-rays." + }, + "9802/astro-ph9802029_arXiv.txt": { + "abstract": "We present long slit spectroscopy for the [OIII] and H$\\alpha$ wavelength ranges along nine different position angles for the Sa Seyfert 1.9 galaxy NGC 2992. Double profiles are present in several regions, suggesting that the gas is not simply following galaxy rotation. A simple kinematical model, which takes into account circular rotation together with a constant radial outflow, seems to be a good approximation to account for the observed kinematics. ", + "introduction": "Disturbed morphologies of ionized nebulosities surrounding active galactic nuclei are frequently observed, mainly in high ionization gas. The morphology is usually interpreted in terms of a conical or biconical shape centered on the nucleus. The spectacular ionizing cone discovered in NGC 5252 by Tadhunter \\& Tsvetanov (1989) is thought to arise from interstellar matter lit up by radiation from the nuclear non stellar continuum escaping the central regions through the hole of an obscuring torus surrounding the nucleus. The situation is far less clear in the other reported cases where the cone may be seen only on one side of the nucleus or has less marked edges (cf Wilson \\& Tsvetanov 1994 for a review). The possibility of having outflows or inflows in cones centered on the nucleus has been invoked to account for the kinematical properties of several Seyferts (e.g. Wilson et al. 1985), but kinematics have been studied in detail only in a few of the Seyferts which show evidence for galactic outflows. In this paper we present new results on the kinematics of the ionized gas in the highly inclined ($i$=70$^\\circ$) Sa Seyfert galaxy NGC 2992, crossed by a disturbed dust lane oriented along the major axis ($\\phi$=15$^\\circ$, RC3 catalogue). This edge-on galaxy is connected by a tidal tail to a close companion, NGC 2993, at a projected distance of 2.9 arcmin (35.5 kpc for H$_0$=50 km s$^{-1}$ Mpc$^{-1}$) and with a 109 km s$^{-1}$ velocity difference (see RC3 catalog), which may well have important perturbing effects on its dynamics. Combined optical broad and narrow band images reveal a complex structure (Durret \\& Bergeron 1987, Wehrle \\& Morris 1988). The [OIII] and H$\\alpha$ images show an arc of emission southeast of the nucleus (which could be interpreted as HII regions in the spiral arm) as well as a finger of emission emerging from the northwest portion of the nucleus, pointing northward. At 20cm radio wavelength, NGC 2992 reveals a radio source of total extent 25 arcsec and major axis PA$\\sim 160^\\circ$, with a one--sided extension along PA$\\sim 130^\\circ$ (Ward et al, 1980; Hummel et al., 1983). At smaller scale an 8-shaped structure, with the nucleus at the crossing-point, is visible at 6cm along PA=160$^\\circ$ and interpreted as limb-brightened bubbles or magnetic arches (Ulvestad \\& Wilson 1984; Wehrle \\& Morris 1988). A comparison of the line emission images with the radio images shows no correlation. Colina et al. (1987) have mapped NGC 2992 by using long slit spectroscopy roughly along the major axis and the two axes given by the high resolution radio map by Ulvestad \\& Wilson (1984). From their kinematical data in the [OIII] lines, they observe blue asymmetric profiles only in the very center of the galaxy. They find that the nuclear and off-nuclear regions are dynamically decoupled, and suggest that there are non-circular motions due to radial flows and tidal interaction with NGC 2993. They interpret their data as being consistent with a radial outflow of gas in a plane which is not coaligned with the galactic plane, rather than outflow within a cone. On the contrary, Tsvetanov et al. (1995) describe NGC 2992 as a good candidate for having large-scale minor axis outflows with velocities up to 200 km s$^{-1}$. Since NGC 2992 is a highly inclined object, it is a good candidate to sample gas motions out of the disk plane. In the following, we present a first order kinematic model of the ionized gas in NGC 2992, based on a set of long slit spectra. ", + "conclusions": "We have obtained the gas kinematics of NGC 2992 by means of long slit data along nine position angles. The results are in general agreement with previous determinations (Heckman et al. 1981, who presented data along PA=120$^\\circ$ with higher spectral resolution but lower spatial resolution; Colina et al. 1987, who presented data of lower spectral and spatial resolution along 3 PAs; Keel 1996, for data with a resolution similar to ours along PA$\\simeq 17^\\circ$). We have modeled the kinematics of NGC 2992 by circular rotation in a gaseous disk to which is added constant radial outflow in the disk plane, as suggested by line splitting. Disk rotation can be accounted for with the following parameters: major axis along $\\phi$=30$^\\circ$, inclination $i$=70$^\\circ$, velocity amplitude $V_0$=250 km s$^{-1}$, and parameters $p$=1.1 and $r_0$=5 arcsec (see Sect. 3). This value of $\\phi$, which allows a fairly good representation of the observed velocity field in the central regions, results to be different from that of the large scale disk as derived by continuum images. This discrepancy could imply that the disk of NGC 2992 is warped, probably due to interaction with NGC 2993. Outflow was modeled along a triangular region of axis PA=120$^\\circ$, with an opening angle of 120$^\\circ$ and a constant velocity of 150 km s$^{-1}$ along the outflow region (measured 5 arcsec from the nucleus). Opening the angle of the outflow to 160$^\\circ$ in the east could give a better fit only for PA=30$^\\circ$ in the northeast, in the region where H$\\alpha$ emission is detected by Wehrle \\& Morris (1988). It can be noted that low excitation gas follows better the pure rotation model in some regions, whereas high excitation gas is better represented when outflow is also included in the model. In our scenario, the outflow takes place radially close to the plane of the gaseous disk, with a spatial extension which is much larger on the east than on the west side. This simple model accounts rather well for the kinematics of NGC 2992, if one excepts the regions located more than 6 arcsec northwest of the nucleus. The kinematics of the gas in that region confirm results by Heckman et al. (1981) and Colina et al. (1987), who reported large velocities and velocity dispersions. As noticed before, these regions with the largest discrepancies may be understood in terms of more complex dynamics and structure, and/or by taking into account the possible line asymmetries due to the dust distribution. Since our spectral resolution is not sufficient to observe line splitting in high excitation lines, and we don't have line ratios, it could be difficult to constrain a more sophisticated model." + }, + "9802/astro-ph9802359_arXiv.txt": { + "abstract": "The interaction of magnetic fields and convection is investigated in the context of the coronal heating problem. We study the motions of photospheric magnetic elements using a time series of high resolution G-band and continuum filtergrams obtained at the Swedish Vacuum Solar Telescope at La Palma. The G-band images show bright points arranged in linear structures (``filigree'') which are located in the lanes between neighboring granule cells. We measure the motions of these bright points using an object tracking technique, and we determine the autocorrelation function describing the temporal variation of the bright point velocity. The correlation time of the velocity is about 100 s. To understand the processes that determine the spatial distribution of the bright points, we perform simulations of horizontal motions of magnetic flux elements in response to solar granulation flows. Models of the granulation flow are derived from the observed granulation intensity images, using a simple 2D model that includes inertia and horizontal temperature gradients; the magnetic flux elements are assumed to be passively advected by this granulation flow. The results suggest that this passive advection model is in reasonable agreement with the observations, indicating that on a time scale of 1 hour the flux tubes are not strongly affected by their anchoring at large depth. Finally, we use potential-field modeling to extrapolate the magnetic and velocity fields to larger height. We find that the velocity in the chromosphere can be locally enhanced at the separatrix surfaces between neighboring flux tubes. The predicted velocities are several km/s, significantly larger than those of the photospheric flux tubes. The implications of these results for coronal heating are discussed. ", + "introduction": "Observations of the Sun in X-rays show that the corona is heated to several million degrees, and that magnetic fields play a key role in this heating process (e.g. Vaiana \\& Rosner 1978; Golub et al. 1980; Pallavicini et al. 1981; Kano \\& Tsuneta 1995, 1996; Shimizu \\& Tsuneta 1997; Yoshida \\& Tsuneta 1996; Falconer et al. 1997). The energy source for this heating must lie in the turbulent convection zone below the photosphere. The interaction of the magnetic field with convective flows produces two types of magnetic disturbances. First, the buffeting of magnetic flux tubes by the granulation flow generates transverse MHD waves (e.g. Steiner, Kn\\\"{o}lker \\& Sch\\\"{u}ssler 1994), which propagate upward along the magnetic flux tubes and dissipate their energy in the chromosphere or corona (see Ofman, Klimchuk \\& Davila 1998, and references therein). Second, in coronal loops the random motions of the footpoints produce twisting and braiding of coronal field lines, which generates field-aligned DC electric currents that can be dissipated resistively (Parker 1972, 1983; Tucker 1973; Rosner et al. 1978; Sturrock \\& Uchida 1981; Heyvaerts \\& Priest 1984; van Ballegooijen 1985, 1986, 1990a; Milano, Gomez \\& Martens 1997). The main difference between these processes is that plasma inertia plays a key role in MHD wave propagation but is not important for the dynamics field-aligned currents. Therefore, magnetic heating mechanisms can be crudely classified as either wave-heating or current-heating mechanisms (for reviews of coronal heating theory, see Narain \\& Ulmschneider 1990, 1996; Zirker 1993). The relationship between coronal heating and the dynamics of photospheric magnetic structures is not well understood. Van Ballegooijen (1985, 1986) proposed that the slow twisting and braiding of coronal field lines by horizontal footpoint motions leads to a {\\it cascade of magnetic energy} within coronal loops. The heating rate predicted by this model is $\\epsilon \\sim B_0^2 D / (4 \\pi L^2)$, where $B_0$ is the coronal field strength, $L$ is the loop length, and $D$ is the ``diffusion constant'' describing the random walk of photospheric flux elements. The latter is given by $D \\sim u_0^2 \\tau_0$, where $u_0$ is the typical velocity of the magnetic elements, and $\\tau_0$ is the velocity coherence time. Observations of the spreading of active regions indicate that $D$ is in the range 150 to 425 $\\rm km^2 ~ s^{-1}$ (DeVore et al. 1985), but such values of $D$ yield coronal heating rates which are a factor $\\sim 40$ too small compared to observed radiative and conductive loss rates (van Ballegooijen 1986). We suggest there may exist short-period motions with velocities of 1-2 km/s which contribute to coronal heating but do {\\it not} contribute to the spreading of active region flux on time scales of days to months (i.e., the motions cannot be described by a ``random walk''). Short-period motions of magnetic elements have indeed been observed (e.g. Berger \\& Title 1996), but the characteristics of such motions are not well understood, and it is unclear whether the velocities are sufficient to explain the observed coronal heating. A number of authors have developed three-dimensional simulation models of the magnetic structure and heating of solar coronal loops (e.g. van Ballegooijen 1988, 1990b; Mikic, Schnack \\& Van Hoven 1989; Longcope \\& Sudan 1994; Hendrix et al. 1996; Galsgaard \\& Nordlund 1996). In these simulations an initially uniform field connecting two parallel plates is considered (the plates represent the photosphere at the two ends of a coronal loop). The footpoints at the boundary plates are subjected to a series of randomly phased, sinusoidal flows patterns. The flows are incompressible, so that the component of magnetic field perpendicular to each boundary plate remains uniform. This is obviously not a good model of the actual conditions in the solar photosphere. Observations indicate that the photospheric magnetic field is highly intermittent, consisting of discrete flux elements with nearly field-free gas in between (e.g. Title, Tarbell \\& Topka 1987; Title et al. 1992). The solar convection plays a key role in producing these flux elements, which are mainly located at the boundaries between granulation cells. This concentration of flux into discrete elements probably has important consequences for coronal heating: if the magnetic field consists of a bundle of topologically distinct flux tubes, then the random motions of these flux tubes in the photosphere will create {\\it tangential discontinuities} at the interfaces between the tubes in the corona (Parker 1972; Glencross 1975, 1980; Rosner et al. 1978; Sturrock \\& Uchida 1981; D\\'{e}moulin \\& Priest 1997). Clearly, to understand where and how the corona is heated, it is necessary to study the dynamics of photosheric flux elements and to develop coronal heating models which take the presence of such flux elements into account. In this paper we focus on the first part of this problem, namely, the dynamics of magnetic flux elements in the photosphere and chromosphere. We analyze observations of G-band bright points obtained at the Swedish Vacuum Solar Telescope (SVST) on La Palma, and derive the autocorrelation function describing the temporal variations of the bright point velocity. We find that the velocity changes on a time scale of about 100 s. We develop an empirical model of the granulation flow and simulate the horizontal motions of flux tubes, assuming they are passively advected by the granulation flow. By comparing the results of these simulations with the observed spatial distributions and velocities of G-band bright points, we show that this passive advection model is in reasonable agreement with the observations. Finally, we use potential-field modeling to extrapolate the magnetic and velocity fields to larger heights (up to 1500 km in the chromosphere). The results indicate that the spreading of the flux tubes with height and their interactions with each other produce plasma flows in the chromosphere with velocities of several km/s, much faster than the velocities of the underlying photospheric flux tubes. This suggests that the coronal heating rate can be significantly enhanced by the three-dimensional geometry of the flux tubes in the photosphere and chromosphere. ", + "conclusions": "Observations of the Sun with high spatial resolution show network bright points (Muller 1983, 1985, 1994; Muller \\& Keil 1983; Muller \\& Roudier 1984, 1992) and ``filigree'' (Dunn \\& Zirker 1973; Mehltretter 1974; Berger et al. 1995), which are small bright features located within the intergranular lanes. The bright points and filigree are seen in the wings of strong spectral lines such as H$\\alpha$ and Ca II H \\& K, in lines formed in the photosphere, and even at continuum wavelengths (with reduced contrast). The widths of these structures is 100 to 200 km, at the limit of resolution of ground-based solar telescopes. The bright points are associated with regions of strong magnetic field (Chapman \\& Sheeley 1968; Title, Tarbell \\& Topka 1987; Simon et al. 1988; Title et al. 1989, 1992; Keller 1992) and correspond to magnetic flux tubes of kilogauss field strength that stand nearly vertically in the solar atmosphere (Stenflo 1973; Stenflo \\& Harvey 1985; Sanchez Almeida \\& Martinez Pillet 1994; see review by Solanki 1993). The granules near network bright points are smaller and more numerous than near a normal intergranular space (Muller, Roudier \\& Hulot 1989). The filigree produce abnormal granulation patterns (Dunn \\& Zirker 1973) and appear to be chains of bright points which fill the intergranular lanes. The dynamical behaviour of bright points has been studied by a number of authors. Muller (1983) found that facular points on the quiet sun are predominantly located in patches at the periphery of supergranule cells, indicating that the magnetic elements are advected by the supergranular flow. The bright points always first appear in the dark spaces {\\it at the junction of several granules}, never inside a granule nor in the space between only two granules. As the granulation pattern evolves, the bright points remain in the intergranular spaces throughout their lifetime, but not necessarily at the junction of several granules like at the time of their first appearance. New bright points have a tendency to appear adjacent to existing points, and 15 \\% of the bright points seem to split into two points which move apart until a separation of 1 to 1.5 arcsec is reached. Muller et al. (1994) measured velocities of 29 isolated bright points and found a mean speed of 1.4 km/s. Strous (1994) studied bright points in a growing active region. Using line-center images taken in Fe I 5576 {\\AA}, he found velocities between 0.26 km/s and 0.62 km/s. Berger \\& Title (1996) measured velocities of 1 to 5 km/s for G-band bright points in the ``moat'' around a sunspot; they showed that the motions are constrained to the intergranular lanes and are primarily driven by the evolution of the granulation pattern. They found that the bright points continually split and merge, with a mean time between splitting events of few hundred seconds. Berger et al. (1998) observed a similar rapid splitting and merging of bright points in an enhanced network region. In the present paper we use observational data from the SVST on La Palma to derive time-dependent granulation flow fields and to simulate the horizontal motions of magnetic elements in the photosphere. The data were collected on 1995 October 5 between 10:57 and 12:08 UT. Observations were made simultaneously with two CCD cameras: one used a 12 {\\AA} bandpass interference filter with a center wavelength of 4305 {\\AA} (G band), and the other used a 54 {\\AA} bandpass filter with a center wavelength of 4686 {\\AA}. Frame selection was based on the G band images: only the three best frames in each 20 s evaluation period were retained for analysis. Both cameras were equipped with phase-diversity beam splitters which put two images on each CCD with a difference in focus position. The images were corrected for seeing effects using Partitioned Phase-Diverse Speckle restoration, and were carefully coaligned using image destretching techniques. The data were space-time filtered to remove the effects of solar p-mode oscillations. The result is a time series of 180 images with very high spatial resolution covering a period of about 70 min. The image scale is 0.083 arcsec/pixel and the mean time between frames $\\Delta t$ = 23.5 s. The field-of-view (FOV) is a 29 by 29 arcsec area near solar disk center containing an enhanced network region. This region is the same as the ``network FOV'' described by L\\\"{o}fdahl et al. (1998), and we refer to their paper for a detailed description the observation and image restoration procedures. The data used here are in essence an early version of these data (L\\\"{o}fdahl 1996). Berger et al. (1998) analyzed ``the network FOV'', where magnetic elements are seen as bright points with high contrast in the G band (which includes the molecular bandhead of CH) and with reduced contrast in the 4686 {\\AA} band (which contains continuum and many absorption lines). The solar granulation shows up with nearly equal contrast in both types of images. The bright points are generally located in the dark intergranular lanes. Subtraction of the G band and 4686 {\\AA} images yields a difference image which shows the bright points and surrounding diffuse emission with unprecedented clarity. Berger et al. identified bright points using a threshold technique applied to these difference images. They define a ``magnetic region'' as the area within their FOV which is covered by bright points at any time during the movie. Using local correlation tracking (LCT) with subfields of 0.4 arcsec, they measure a mean granulation flow velocity of 0.641 km/s inside this magnetic region and 0.997 km/s in the surrounding quiet region. They also followed the motions of bright points using an object tracking technique, and found that the bright points have a broad velocity distribution which peaks at 0.1 km/s but extends to several km/s; the mean velocity is 0.815 km/s. The bright points continually split up and merge; the average time between merging and splitting events is 220 s. Some objects can be followed for the entire 70 min duration of the movie. Following Berger et al. (1998), we construct a time series of ``magnetic'' difference images by subtracting the 4686 {\\AA} continuum images from the corresponding G-band images: \\begin{equation} I_{magn} (x,y,t) = I_{4305} (x,y,t) - I_{4686} (x,y,t) , \\end{equation} where $x$ and $y$ are horizontal coordinates on the Sun, $t$ is the time, and $I_4305$ and $I_4686$ are intensities normalized to the mean intensity in non-magnetic areas of the frame. Examples of such difference images are shown in Figures \\ref{magn_image}a and \\ref{magn_image}b (frames 40 and 177 of the time sequence). Note that the bright points form linear structures which fill the intergranular lanes and sometimes surround the granules on all sides. We also construct a time series of magnetic ``masks'' by smoothing the $I_{magn}$ images in space and time (running average over 5 frames) and applying a $2 \\sigma$ threshold. These masks define the general areas where the bright points are located. \\placefigure{magn_image} The simulations in \\S 4 show that the observed velocities and spatial distribution of the G-band bright points are consistent with passive advection by the granulation flow, i.e., we find no evidence for deep-seated flows different from the observed granulation pattern. To reproduce the observed ``filigree'' structures, we had to take into account that the photospheric flux elements are nearly incompressible ($B \\sim 1500$ G) and that the intergranular lanes are very narrow ($w \\sim 100$ km). In contrast, in models with broader lanes most of the magnetic elements collect in a few large clusters at the vertices where several lanes intersect. Therefore, the observed filigree structures provide indirect evidence for the existence of narrow intergranular lanes with widths $\\sim 100$ km. Note that the present granulation flow model is based on the {\\it observed} granulation pattern and therefore includes the effects of the magnetic field on the granulation flow (longer lifetime of granules; reduced cell size). Three-dimensional models of solar convection (e.g. Stein \\& Nordlund 1994) show the presence of small-scale vortical flows within the intergranular lanes on scales $\\sim 100$ km. Such small-scale flows likely play an important role in breaking up the ``filigree'' into distinct magnetic elements, as is observed in high-resolution G-band images. Indeed, bright points sometimes exhibit rotational motions about each other (Berger \\& Title 1996). In the present paper we neglected the vorticity of the granulation flow and its effects on the dynamics of flux tubes. The reason is that the observational data do not provide information about the vorticity of granulation flows on small spatial scales. Using correlation tracking, it is possible to measure vorticity on length scales of several granules (e.g. Wang et al. 1995), but such large-scale vorticity probably has little effect on the small-scale structure and dynamics of magnetic elements within the intergranular lanes. Clearly, future modeling of flux tube dynamics will need to take vortical flows into account, and work along these lines is in progress (Nordlund \\& Stein, in preparation). To measure small-scale vortical flows directly (other than by tracking magnetic elements) will require observations with much higher spatial resolution than are presently available. How effective are the observed motion in generating MHD waves? In \\S 3 we show that the autocorrelation function of the bright point velocity is approximately given by: \\begin{equation} C( \\tau ) \\equiv < v_x (t) v_x (t+ \\tau ) > = \\frac{\\sigma^2} {1 + ( \\tau / \\tau_0 )^2 } , \\end{equation} where $\\sigma^2 \\approx 0.12 ~ {\\rm km^2 s^{-2}}$ and $\\tau_0 \\approx 100$ s. The velocity power spectrum is obtained by taking the Fourier transform of the autocorrelation function: \\begin{equation} P( \\omega ) \\equiv \\frac{1}{2 \\pi} \\int_{- \\infty}^{+ \\infty} C( \\tau ) e^{i \\omega \\tau} d \\tau = \\frac{1}{2} \\sigma^2 \\tau_0 e^{- \\tau_0 | \\omega |} . \\end{equation} Hence, most of the power is contained in frequencies $\\omega < \\tau_0^{-1}$ $\\approx 0.01$ rad/s. This is significantly less than the acoustic cutoff frequency in the photosphere, which is given by $\\omega_{ac}$ = $c/(2H) \\approx 0.03$ rad/s, where $c \\approx 7$ km/s is the sound speed and $H \\approx 115$ km is the pressure scale height. Therefore, it is doubtful that the observed motions are very effective in generating longitudinal tubes waves, which are evanescent at these frequencies. A number of authors have proposed that the corona is heated by {\\it transverse} tube waves (Hollweg 1984; Choudhuri, Auffret \\& Priest 1993; Choudhuri, Dikpati \\& Banerjee 1993). Muller {\\it et al.~}(1994) measure a mean speed of network bright points of 1.4 km/s and conclude that there is sufficient energy in these motions to heat the quiet corona. However, it is not clear that the observed velocity is entirely due to transverse waves: part of the mean velocity must be due to long-period motions which are not effective in generating transverse tube waves. The cutoff frequency for tranverse tube waves is given by \\begin{equation} \\omega_c^2 = \\frac{g}{8H} \\frac{1}{2 \\beta +1} , \\end{equation} where $g$ is the acceleration of gravity, $\\beta = 8 \\pi p_i /B^2$ is the ratio of gas pressure and magnetic pressure, and the thin flux tube approximation has been used (Spruit 1981). Using $\\beta$ = 0.3, we find $\\omega_c = 0.0135$ rad/s. If we assume that only the high frequency modes with $| \\omega | > \\omega_c$ can generate transverse waves, then the amplitude of these waves is given by: \\begin{equation} u_t^2 \\equiv \\sigma^2 \\tau_0 \\int_{\\omega_c}^{\\infty} e^{- \\tau_0 \\omega} d \\omega = \\sigma^2 e^{- \\tau_0 \\omega_c } \\approx 0.031 ~~~ [ {\\rm km^2 s^{-2}} ] , \\end{equation} which yields $u_t \\approx 0.18$ km/s. The energy flux in such waves is given by $F = \\rho u_t^2 v_A$, where $\\rho$ is the mass density, and $v_A \\equiv B/ \\sqrt{4 \\pi \\rho}$ is the Alfv\\'{e}n speed. Using $\\rho = 7 \\times 10^{-8}$ $\\rm gr/cm^3$ for the density within the tubes in the low photosphere, we find that the energy transported by these waves per unit magnetic flux is $F/B$ = $u_t^2 \\sqrt{ \\rho / ( 4 \\pi )}$ = $2.3 \\times 10^4$ $\\rm erg / cm^2 / s / G$. Hence, in plage regions with magnetic flux density $\\sim 50$ G, we expect a mean energy flux $\\sim 10^6$ $\\rm erg / cm^2 / s$. This is comparable to the observed radiative and conductive losses from the corona above plage regions (Withbroe and Noyes 1977). Extrapolations of magnetic and velocity fields to larger heights (see \\S 5) show that the chromospheric velocity can be locally much larger than the horizontal velocities of the flux tubes in the photosphere. This enhancement is due to the 3D geometry of the flux tubes and the dynamics of field lines near magnetic null points where different flux tubes interact. This result may have important implications for models of coronal heating. First, the coronal heating rate varies as the {\\it square} of the footpoint velocity, hence the velocity enhancement found here could significantly increase the coronal heating rate. Second, the enhanced velocities are localized near separatrix surfaces, and therefore could affect the way in which electric currents build up at such surfaces (Glencross 1975, 1980; D\\'{e}moulin \\& Priest 1997). Finally, the predicted chromospheric velocities are so large that plasma inertia cannot be neglected; quasi-static potential-field models cannot accurately describe the flows that occur where neighboring flux tubes interact. Clearly, to understand how magnetic energy is stored and dissipated in the corona, a more realistic {\\it dynamical} model of interacting flux tubes is needed." + }, + "9802/astro-ph9802271_arXiv.txt": { + "abstract": "Models of relativistic jets filled with ultrarelativistic pair plasma are very successful in explaining the broadband radiation of $\\gamma$-ray blazars. Assuming that the initial injection and cooling of ultrarelativistic pair plasma in an AGN jet has occurred, producing the observed high-energy $\\gamma$-ray radiation, we investigate the further evolution of the pair plasma as it continues to move out from the central engine. The effects of thermalization and reacceleration, the emission of pair bremsstrahlung and annihilation radiation and the bulk Compton process, and the possible application to MeV blazars are discussed. A model calculation to the special case of PKS~0208-512 is presented. ", + "introduction": "The detection of high-energy $\\gamma$-ray emission from more than 60 blazars with EGRET is a challenge and at the same time a constraint of fundamental importance for emission models (von Montigny et al. 1995). A large fraction of these blazars exhibits variability at $\\gamma$-ray energies on time scales of days to months (Mukherjee et al. 1997). The optical counterparts of the majority of EGRET detected AGN are known as BL Lacertae objects and optically violent variable quasars (OVV). At radio wavelengths, all blazars can be recognized as bright, compact sources with a flat synchrotron spectrum emanating from outflowing plasma jets that are nearly aligned with our line-of-sight. Relativistic beaming is required in the objects in view of the luminosity and variability time scales (Dermer and Gehrels 1995), in accord with VLBI observations indicating that superluminal motion is a common feature in this class of AGN (e.g. Wehrle et al. 1994, Pohl et al. 1995, Barthel et al. 1995, Krichbaum et al. 1995). The strongest EGRET blazar detections can be characterized by a single power-law spectrum with differential photon spectral indices between $\\alpha = 1.5$ and $\\alpha = 2.7$ (Thompson et al. 1995). For individual sources the spectral index is correlated with the flux level, and there may also be deviations from the power-law behaviour both below 70 MeV and above a few GeV (Pohl et al. 1997). The combined OSSE/COMPTEL/EGRET measurements generally indicate spectral breaks at a few MeV (Williams et al. 1995, McNaron-Brown et al. 1995) in the sense that the spectra below 1 MeV are harder than in the EGRET range. At medium $\\gamma$-ray energies observable by COMPTEL, PKS 0208-512 has been identified as an AGN with flaring properties at MeV energies (Blom et al. 1995). There is now evidence that it may belong to a class of `MeV-blazars' that are occasionally exceptionally bright MeV sources (see Bloemen et al. 1995). Since in these objects the bright emission appears to be confined to a relatively narrow energy range, the discussion has focused on models involving a broad blue-shifted $e^+/e^-$ annihilation line that is Doppler boosted in a relativistic jet (Roland and Hermsen 1995). If this interpretation applies, MeV-blazars provide a unique tool to study astrophysical particle beams (e.g. Schlickeiser 1996). Furthermore, MeV-blazars may confuse the analysis of galactic sources of $\\gamma$-ray line emission (Pohl 1996). In an earlier paper (B\\\"ottcher, Mause \\& Schlickeiser 1997; hereafter BMS) we have investigated the temporal evolution of ultrarelativistic pair plasmas in jets of quasars and BL-Lac objects and have demonstrated that their broadband spectra can well be explained as the resulting synchrotron and inverse-Compton radiation from a cooling ultrarelativistic nonthermal pair plasma in a relativistic jet. A decisive parameter for the evolution of single ultrarelativistic plasma components inside an AGN jet is the density of pairs injected into the jet. Broadband fits, covering the radio to $\\gamma$-ray regime of the electromagnetic spectrum, to blazars generally require particle densities of order $n_e \\ukl 10^3$~cm$^{-3}$. Nevertheless, fits to different objects suggest that the value of the pair density in relativistic jets ejected by active galactic nuclei varies over several orders of magnitude. Jets of very high density ($n \\ugr 10^{5}$~cm$^{-3}$) can also produce the broadband spectra at least of $\\gamma$-ray active flat-spectrum radio quasars whose bolometric luminosity is clearly dominated by the $\\gamma$-ray emission. In this paper, we argue that after an initial phase of rapid cooling, governed by synchrotron and inverse-Compton energy losses, the relativistic pair plasma inside such components is likely to attain a quasi-thermal distribution. If the jet remains well collimated and cooling (e. g. through adiabatic losses) remains very efficient through the transrelativistic phase, pair bremsstrahlung and pair annihilation become efficient. The resulting radiation spectrum peaks around several MeV, which has been suggested previously to be responsible for the observed MeV bump in MeV blazars (Henri et al. 1993, Roland \\& Hermsen 1995, B\\\"ottcher \\& Schlickeiser 1996). However, while generally the high-energy radiation from a few ultrarelativistic jet components, ejected over the typical EGRET variability time scale, is appropriate to model EGRET spectra, the pair annihilation and bremsstrahlung radiation from only a few components after cooling and thermalization is usually too weak to explain the MeV blazar phenomenon. Therefore, a quasi-continuous supply of mildly relativistic pair plasma into the jet is required in order to fit the MeV bump in MeV blazars with pair annihilation radiation. Alternatively, if very close to the central accretion disk the inverse-Compton cooling rate is balanced by reacceleration by hydromagnetic turbulences the pair plasma in the jet will thermalize at relativistic temperatures and will continuously Compton upscatter external radiation from the accretion disk very efficiently, producing another bump at keV --- MeV energies by the so-called bulk Compton process (Sikora et al. 1997). We know from the EGRET data of many blazars that the GeV emission is not that of a single injection of particles, but is more likely to result from more or less regularly repeating injection events with varying energy input into the ejected particles. A second important point is the correlation between flux and spectral index in the EGRET range. The spectral hardening during outbursts indicates a more efficient acceleration of particles. Even for highly variable sources like PKS~0208-512, PKS~0528+134 (Collmar et al. 1997) and others the low-energy $\\gamma$-ray continuum varies with small amplitudes and not in phase with the variability in the EGRET range. To be more realistic we may thus either assume that injections occur in regular time intervals, so that they would not influence each other in their evolution, or that injection occurs in a quasi-steady manner, such that relativistic pair plasma is continuously injected. For ease of both computing and exposition we will not consider the effect of the finite light travel time within the volume occupied by pair plasma, which limits our predictions of variability to time scales longer than the light travel time through individual plasma blobs. In section 2, we give a short overview of the different elementary processes which play a role for the evolution of a relativistic pair plasma component of an AGN jet after the initial phase of rapid cooling and which are usually not considered in models of ultrarelativistic jets. In section 3, we describe numerical simulations and an analytical approximation, applicable under special conditions, to follow the evolution of the pair plasma through the transrelativistic phase. We find that, depending basically on the particle density and the magnetic field, there are two different ways how a quasi-thermal distribution can be established. The relevant physics of these quasi-thermal plasmas is described in section 4. Section 5 contains a model calculation for the typical MeV blazar PKS~0208-512. We summarize in section 6. ", + "conclusions": "Motivated by the standard model for $\\gamma$-ray emission from blazars invoking jets filled with ultrarelativistic pair plasma oriented at a small angle with respect to the line of sight, we investigated the further evolution of the leptonic material inside such a jet after the initial phase in which the high-energy ($> 100$~MeV) $\\gamma$-ray radiation is produced. We demonstrated that either the action of elastic scattering of particles off each other or the balance of reacceleration by turbulent plasma waves to radiative losses can, under a wide range of parameters, establish a quasi-thermal particle distribution inside the jet, irrespective of the initial pair distribution at the time of injection. Extreme conditions (high density, very weak magnetic field, injection close to the accretion disk) are necessary to achieve mildly relativistic temperatures ($\\Theta \\sim 1$). Radiation from a quasi-thermal pair plasma can produce the temporary MeV bumps observed in MeV blazars. This can be accomplished most easily via inverse-Compton scattering of external soft radiation by a quasi-continuously filled jet. Alternatively, the MeV bump could be produced by pair annihilation radiation, if cooling of the pair distribution is very efficient and the dilution of the plasma through jet expansion is negligible. We found that only the unrealistic assumption of a cylindrical jet yields an acceptable fit to the observed MeV bump, if the pair plasma producing this bump is subject to the same processes (in the same environment) as the ultrarelativistic pair plasma producing the EGRET spectrum. In this picture, an MeV outburst reflects an increasing supply of relativistic electrons and positrons into the jet on much longer timescales than those typically observed in EGRET outbursts at higher energies. It is only weakly or not at all correlated to activity in the EGRET range which corresponds to an increased energy density at the acceleration site, leading to a harder particle spectrum of the injected pairs. If the bulk Compton process is the dominant mechanism for the production of the MeV bump in MeV blazars, then there is no obvious reason for a ``universality'' of the bump photon energy at 3 --- 10~MeV. Mainly depending on the magnetic field strength, the level of hydromagnetic turbulence and the injection height of the plasma blob above the accretion disk, similar sources with bulk-Compton radiation bumps at different energies should exist as well. In general, our treatment is also applicable to the so-called hadronic jet models, where ultrarelativistic protons initiate a pair cascade in the jet (Mannheim et al. 1991, Mannheim \\& Biermann 1992, Mannheim 1993). However, although the jet simulation code of BMS is able to handle problems in which the blob is optically thick to $\\gamma\\gamma$ pair production, it is clearly beyond the scope of this paper to follow the proton-initiated pair cascades in detail. Nevertheless, a few general conclusions may be drawn from the fact that in those models, the magnetic fields are fairly high and a considerable level of turbulence might be present. Our simulations indicate that at magnetic fields of $B \\gtrsim 10$~G, as usually assumed in hadronic jet models, and the moderate pair densities expected to result in the cascades, the pair plasma will maintain a highly relativistic temperature, $\\Theta \\gtrsim 100$. Thus, no pair annihilation feature will be observable; instead, strong synchrotron and SSC bumps at $\\sim 3 \\cdot 10^7 \\cdot D/(1+z) \\, (B/G) \\, \\Theta^2$~Hz and $\\sim 4 \\cdot 10^8 \\cdot D/(1+z) \\, (B/G) \\, \\Theta^4$~Hz, respectively, would be expected from the thermalized cascade plasma." + }, + "9802/astro-ph9802047_arXiv.txt": { + "abstract": "We have studied the optical counterparts of X1755$-$338 and X1658$-$298 in their X-ray ``off'' states. The first observations of X1755$-$338 in quiescence show that the counterpart, V4134 Sgr, has faded by more than 3.5~mag in $V$. If the mass donor in the system is an M0V star as implied by the period, our upper limits on the brightness of the counterpart suggest that it is more distant than 4~kpc. We observed V2134 Oph, the optical counterpart of X1658$-$298, on several occasions in April/May 1997 and found the source only $\\sim$~1~mag fainter than when it is X-ray bright. Contemporaneous X-ray data confirm that the source remains in the quiescent state during our optical observations. Our optical lightcurve, folded on the 7.1~hour orbital period, does not show any modulation across the binary cycle. It is possible that the absence of detectable X-ray emission, despite the indication for an accretion disk and activity in the system, is related to structure in the disk that permanently obscures the central X-ray source. The optical properties of V2134~Oph are unique among the known X-ray transients. ", + "introduction": "About ten systems among the low mass X-ray binaries (LMXBs) exhibit irregularly-shaped, recurrent dips in their X-ray lightcurves. These X-ray dippers are believed to be high inclination systems in which azimuthal accretion disk structure extends above the plane of the binary and periodically blocks the line of sight to the central compact object. Modeling shows that for the majority of systems this structure must be at the disk edge, at the impact point of the accretion stream (White, Nagase \\& Parmar 1995 and references therein). The recurrence time of the dips is assumed to reflect the orbital period. The optical/IR emission from dippers also varies on the orbital period, and originates largely from reprocessing of X-rays in the outer regions of the system. Observations of LMXBs during X-ray quiescence offer the opportunity to study the optical counterpart ``uncontaminated'' by contributions to the optical/IR from reprocessed X-radiation in the disk and secondary. Light from the accretion disk usually dominates in the X-ray active state. During quiescence, the mass donor becomes visible which allows the determination of its spectral type and radial velocity curve. Radial velocity measurements can provide the mass function and scale of the system and help to distinguish whether the compact object is a neutron star (NS) or black hole (BH). ", + "conclusions": "\\subsection{X1755$-$338} X1755$-$338 was first noted for its very soft X-ray spectrum by Jones (1977). Observations with {\\it Einstein} identified it as a BH candidate due to its location in the ultrasoft region of the X-ray color-color diagram (White \\& Marshall 1984). The {\\it Einstein} spectra also indicated a lower than expected column density for a source close to the Galactic Bulge. {\\it EXOSAT} observations by White et al.\\ (1984) revealed recurrent dips of 30~min duration with a period of 4.4~hours. Unlike in other X-ray dippers, the hardness ratio did not change between dip and non--dip periods. One explanation offered for this energy independence of the X-ray dips was a reduction in the metallicity of the absorbing medium by a factor of 600 from cosmic abundance values. Church \\& Balucinska--Church (1993) fitted the same observations with a two-component model, and were able to reproduce the energy-independence of the dips with absorption in material of cosmic abundances. They interpret the blackbody component of their model as emission from the boundary layer between the accretion disk and the surface of a NS, calling into question the BH candidacy of X1755$-$338. Pan et al.\\ (1995) reported the detection of a hard power-law tail in addition to an ultrasoft component in the X-ray spectrum of X1755$-$338 and strongly argued in favor of X1755$-$338 being a BH candidate. Simultaneous {\\it Ginga} and {\\it Rosat} observations confirm the hard tail and also show the iron 6.7~keV line indicating that the accreted material is {\\it not} extremely metal-deficient (Seon et al.\\ 1995). A faint, blue star with a featureless spectrum was suggested as the optical counterpart of X1755$-$338 by McClintock et al.\\ (1978). Mason, Parmar \\& White (1985), hereafter MPW, obtained simultaneous optical and {\\it EXOSAT} observations of X1755$-$338. Their detection of sinusoidal modulation in the $V$ band with an amplitude of 0.4~mag and a period of 4.46~hours confirmed the identification of the counterpart (V4134 Sgr). The optical minimum occurs 0.15 cycles later than the center of the X-ray dip, consistent with the idea that the X-rays are being absorbed in material located at the impact point of the accretion stream with the accretion disk. In January 1996, X1755$-$338 was observed by the Rossi X-ray Timing Explorer (RXTE) in an X-ray off--state (a factor of 100 fainter than usual) for the first time (Roberts et al.\\ 1996). We obtained optical observations in order to determine the nature of the mass donating star. The RXTE All Sky Monitor (ASM) lightcurve (available on the Web) confirms that the source remained in the off--state during our observations. Figure~\\ref{f-fc1} displays our $V$ band observation of the X1755$-$338 field from 1997 April 28 UT. For comparison, a section of the finding chart which shows the source in its bright state is also reproduced (with the kind permission of K. Mason). The counterpart has faded from its on--state brightness of $V=18.5$ (MPW) to undetectable magnitudes. We obtained upper limits on the counterpart brightness of $V>22$, $R>21.5$ and $I>21$. Our $V$ magnitudes for the comparisons 1 through 4 (see Table~\\ref{t1}) agree with the measurements of MPW within the uncertainties. The distance to X1755$-$338 is not well determined. Estimates in the literature range from 1 to 9~kpc and values for the visual extinction lie between $1< A_V < 2$ (MPW). \\begin{figure}[ht] \\epsscale{0.8} \\plotone{fig1.ps} \\caption[]{{\\it Left:} $V$ band exposure of V4134~Sgr (X1755$-$338), marked X, during outburst in 1984 (MPW). North is up and east is to the left. {\\it Right:} Our 900~s $V$ band exposure taken on 1997 April 28 UT with the CTIO 1.5~m telescope of the same field. Notice that the counterpart of X1755$-$338 has vanished. \\label{f-fc1}} \\end{figure} In order to derive a distance from our magnitude limits, we must consider the origin of the optical quiescent emission in X1755$-$338. In quiescence, the main light source in the system is most likely the mass donating star. There is evidence for a disk even in quiescence in soft X-ray transients, but it generally does not contribute much of the light in the system (McClintock \\& Remillard 1990; Marsh, Robinson \\& Wood 1994). X1755$-$338 has been X-ray bright for over 20~years, which is very different from typical transient behavior (relatively short X-ray outbursts separated by long periods of quiescence). For the remainder of the discussion we assume that all of the light originates from the mass donor. Using the period--mass relations of Frank, King \\& Raine (1992) and Warner (1995) together with the 4.46~h period of the system (MPW) results in mass estimates for the donor star of 0.49~M$_{\\sun}$ and 0.42~M$_{\\sun}$, respectively. If the stars are normal main sequence stars, these masses imply a spectral type of K9-M1 (Allen 1973). Our upper limits in $I$ require a M0V star with $A_V=1$ to be at a distance d$>5$~kpc and at d$>4$~kpc for $A_V=2$ (using the extinction relationship of Cardelli, Clayton \\& Mathis 1989). Alternatively, if we make no assumption about the spectral type of the donor star, at a distance of 9~kpc with $A_V=2$ the spectral type of the secondary would have to be later than K3 in order not to be detected in $I$. \\begin{table} \\dummytable\\label{t1} \\end{table} \\subsection{X1658$-$298} X1658$-$298 is a transient X-ray burst source discovered in 1976 by Lewin, Hoffmann \\& Doty (1976). Observations during a temporary brightening of the source in 1978 showed dips in the X-ray lightcurve. Detailed analysis of the combined 1976--1978 data set by Cominsky \\& Wood (1984, 1989) revealed that some of the dips are in fact eclipses of the central X-ray source by the mass donating star with a recurrence period of 7.1~hours. The dipping activity lasts for about 25\\% of the orbital cycle followed by an eclipse of $\\sim 15$~min duration. X1658$-$298 entered an X-ray off--state in 1979 and has not been detected in X-rays since. The optical counterpart was identified during the 1978 X-ray outburst with a faint ($V=18.3$), blue star (V2134~Oph) by Doxsey et al.\\ (1979). Spectroscopic observations showed a blue continuum with emission lines of \\ion{He}{2} $\\lambda$4686 and the \\ion{C}{3}/\\ion{N}{3} $\\lambda$4640/4650 blend (Canizares, McClintock \\& Grindlay 1979). In 1979 June, the counterpart was detected with $V=21.5$; about a month later it was undetectable with a magnitude limit of $V>23$ (Cominsky, Ossmann \\& Lewin 1983). There is evidence that the source has been brightening gradually since then. Cowley, Hutchings \\& Crampton (1988) mention obtaining a spectrum of the ``very faint source'' in 1986. Shahbaz et al.\\ (1996) display a featureless spectrum of V2134~Oph taken in 1988 and estimate $V=20.7$ for the brightness of the counterpart. Navarro (1996) reported that spectra taken in 1992/1993 showed H$\\alpha$ in emission, but that simultaneous {\\it VLA} and {\\it Ginga} observations failed to detect any radio or X-ray flux. We observed V2134~Oph on several occasions in April/May 1997 and found the source at a mean brightness of $V=19.54$, only $\\sim$~1~mag fainter than when it is X-ray active. Since the only finding chart of X1658$-$298 was taken on photographic plates and is of limited quality, one of our $I$ frames is displayed in Figure~\\ref{f-fc2}. We performed astrometry on our CCD frame to ensure that the observed star is indeed V2134~Oph and not a close companion unrelated to the X-ray source. Our position for V2134~Oph is 17:02:06.42, $-$29:56:44.33 (J2000) with an internal uncertainty of 0.1\\arcsec\\ in each coordinate. Doxsey et al.\\ (1979) list 17:02:06.37, $-$29:56:43.23 for the counterpart (no error estimate is given), which only differs by 1\\arcsec\\ in declination. It therefore appears that the correct counterpart has been observed. A comparison between our $B-V$ color ($B-V=0.80$) and that of Doxsey et al.\\ (1979, $B-V=0.37$) shows that the source is presently substantially redder than during outburst. \\begin{figure}[ht] \\epsscale{0.45} \\plotone{fig2.ps} \\caption[]{600~s $I$ band exposure of the X1658$-$298 field. The field size is 96\\arcsec$\\times$96\\arcsec. North is up and east is to the left. V2134~Oph, the counterpart of the X-ray source, is flagged. \\label{f-fc2}} \\end{figure} Cominsky \\& Wood (1984) discuss the properties of possible secondary stars consistent with producing a 15~min eclipse in a 7.1~hour binary. The limiting cases are a 0.3~M$_{\\sun}$ (M5) star that does not fill its Roche lobe viewed at an inclination of $i=90^\\circ$ and a 0.9~M$_{\\sun}$ (G5) Roche lobe filling star viewed at an inclination of $i=71.5^\\circ$. Using the period--mass relations of Warner (1995) and Frank, King \\& Raine (1992) results in mass estimates of 0.75--0.78~M$_{\\sun}$ for the donor star, which corresponds to a K0 main sequence star. Cominsky (1981) derived a distance of 15~kpc for X1658$-$298. At that distance, the small reddening ($E_{B-V}=0.3$) and the limit of $V>23$ imply that the secondary's spectral type is later than K2. A K star companion is very typical for a large amplitude X-ray transient. If we assume a K3 star mass donor at 15~kpc it would contribute less than 5\\% to the current system brightness, so that most of the light is due to an accretion disk. This is supported by our dereddened colors, which are marginally consistent (within the uncertainties) only with an F5 spectral type. Such a spectral type is too luminous for the previously observed faint source states. Our optical lightcurve, folded on the 7.1~hour orbital period, does not show any modulation across the binary cycle (Figure~\\ref{f-lc}). This is somewhat surprising since many transients in quiescence display photometric variability due to ellipsoidal variations. However, these modulations could be masked by the presence of the apparently luminous accretion disk in X1658$-$298. An accretion disk with non--uniform illumination (e.g. with a hot spot) would also cause optical variability when viewed across the binary orbit. We place an upper limit of 0.02~mag (99\\% confidence) on the semi--amplitude of any sinusoidal modulation in the folded lightcurve. Unfortunately we cannot compare the outburst and quiescent variability characteristics, since there are no data available on the presence or absence of optical modulation during the 1978 outburst. We do not observe any eclipses, but due to our limited phase coverage and the short duration of the eclipses they could have been easily missed. \\begin{figure}[ht] \\epsscale{0.80} \\plotone{fig3.ps} \\caption[]{$I$ band lightcurve of V2134~Oph, the optical counterpart to X1658$-$298, and comparison star 1 (see Figure~\\ref{f-fc2}) folded on the 7.1~hour period of X1658$-$298. The 1$\\sigma$ error bars have been derived from the scatter in the comparison lightcurves. \\label{f-lc}} \\end{figure} The optical behavior of X1658$-$298 is very different from other X-ray transients, which are characterized by a steep increase in optical brightness during an X-ray outburst and roughly constant magnitudes in quiescence. X1658$-$298, in contrast, seems to be brightening gradually in the optical with no accompanying evidence for an increase in the X-ray emission. The RXTE ASM confirms that the source remains at $< 1$ mCrab during our observations. It is possible that the absence of X-ray emission, despite the indication for an accretion disk and activity in the system, is related to the apparent high inclination of X1658$-$298. Changes in the accretion rate could have altered the structure in the outer parts of the accretion disk in such a way that it now permanently obscures the X-rays emanating from the central source. This would also block our view of the inner, hotter regions of the disk which is consistent with the redder colors of our observations." + }, + "9802/astro-ph9802337_arXiv.txt": { + "abstract": "Accurate synthetic models of stellar populations are constructed and used in evolutionary models of stellar populations in forming galaxies. Following their formation, the late type galaxies are assumed to follow the Schmidt law for star formation, while early type galaxies are normalized to the present-day fundamental plane relations assumed to mimic the metallicity variations along their luminosity sequence. The stars in disks of galaxies are distributed with the Scalo IMF and in spheroids with the Salpeter IMF. We show that these assumptions reproduce extremely well the recent observations for the evolution of the rate of star formation with redshift. We then compute predictions of these models for the observational data at early epochs for various cosmological parameters $\\Omega, \\Omega_\\Lambda$ and $H_0$. We find good match to the metallicity data from the damped $L_\\alpha$ systems and the evolution of the luminosity density out to $z\\simeq 1$. Likewise, our models provide good fits for low values of $\\Omega$ to the deep number counts of galaxies in all bands where data is available; this is done without assuming existence of extra populations of galaxies at high $z$. Our models also match the data on the redshift distribution of galaxy counts in $B$ and $K$ bands. They also provide good fits to the observed colors. We compute the predicted mean levels and angular distribution of the cosmic infrared background produced from the early evolution of galaxies. The predicted fluxes and fluctuations are still below the current observational limits, but not by a large factor. Finally, we find that the recent detection of the diffuse extragalactic light in the visible bands requires for our models high redshift of galaxy formation, $z_f \\geq$(3-4); otherwise the produced flux of the extragalactic light at optical bands exceeds the current observational limits. ", + "introduction": "The epoch and process of galaxy formation are still a matter of considerable debate despite substantial recent observational and theoretical progress. On the observational side, it is becoming increasingly clear that galaxies must have formed early on in the evolution of the Universe with the farthest galaxies known to date having redshifts of 5 and beyond (\\pcite{Franx+97}; \\pcite{Dey98}; \\pcite{Hu98}). Similarly, such high redshift of formation of the first stellar populations in galaxies is indicated by the existence of galaxies at moderately high redshifts, $z \\sim 1.5$, but which contain old stellar populations of about 3.5-4 Gyr (\\pcite{Dunlop+96}, \\pcite{Dunlop_98}). On theoretical front, the existence of galaxies at these redshifts coupled with the data on the present day large-scale galaxy distribution allow one to reconstruct the spectrum of the pregalactic density field (\\pcite{Kash_98}) independently of any assumed cosmological paradigms. Other data do not involve record numbers for redshifts of still only a handful of galaxies, but are just as important. Such data come from the deep galaxy counts probing galaxies in various spectroscopic bands, from blue to near-infrared. It too constrains both galaxy evolution and global cosmological parameters. The data on galaxy counts are now coupled with the newly obtained measurements of the luminosity density produced by galaxies in UV, V and J (1.25 micron) bands at $z\\leq 1$ from \\scite{Lilly+96} and the new data on the evolution of cosmic abundances out to high redshift (\\pcite{Pettini+97}). Furthermore, there are now independent measurements of the star formation rate out to high redshifts (\\pcite{Madau+96}); these show that star formation increases out to $z \\simeq 2$ with a possible peak at $z \\sim$ (2-3). All such data allow one to reconstruct the early evolution of galaxies and stellar populations in the Universe thereby providing an important test of galaxy formation processes and the underlying cosmology. In computing such evolution one must necessarily normalize galaxy populations to the present-day data, e.g. galaxy luminosity function in the relevant bands and the morphological mixes. The present-day luminosity function of galaxies has now been measured accurately in B (\\pcite{Loveday+92}) and K bands (\\pcite{Gardner+97}) and the morphological mixes at the present epoch are also well determined (\\pcite{Marzke+94}). With the input of the Initial Mass Function (IMF) for the various galaxy types and specifying their star formation history, one can uniquely compute individual galaxy populations out to very early times. Individual galaxy observations, however, provide only a limited amount of information on the overall evolution of galaxies and the Universe, and are expensive in terms of time involved and area covered. On the other hand, diffuse background radiation fields left over from galaxy formation and evolution contain cumulative information about the entire evolution of the Universe including radiation from objects inaccessible to telescopic studies. Cosmic Infrared Background (CIB) occupies a unique space among the various diffuse background produced by galaxies. The reason is that the bulk of any stellar light emitted at early times will reach the observer shifted into the near to far-IR. CIB thus probes early galaxy formation and evolution and contains cumulative information on the history of the Universe over redshifts between the epoch of the last scattering surface, probed by the microwave background, and $z\\sim 0$, probed by the surveys in visible bands. On the observational side, no detection of the putative CIB has been made at wavelengths below 100 $\\mu$m. At wavelengths beyond 150 $\\mu$m there are claims of possible detections from the COBE DIRBE (\\pcite{Hauser+98}; \\pcite{Schelegel+98}) and FIRAS (\\pcite{Puget+96}) maps. In the near-IR, arguments based on chemical evolution predict levels of CIB around $\\sim 10$ \\nwm2sr (\\pcite{Stecker+77}). The Galactic and zodiacal foregrounds are very bright at these wavelengths so it is very difficult to reach such limits directly as analysis of d.c. levels of DIRBE maps indicates (\\pcite{Hauser+98}). On the other hand, comparable levels can be reached with fluctuations analysis of the near-IR CIB (\\pcite{Kash1+96}; \\pcite{Kash+96}; \\pcite{KMO98}) and there is hope of reaching the CIB levels directly with this method applied to other surveys. In optical bands, \\scite{Vogeley_98} has derived strong limits of the diffuse background at $R$ and $B$ from the fluctuations analysis of the Hubble Deep Field. At the same time, significant progress has been made in the last few years in theoretical understanding of evolution of stellar populations as new opacities have become available (OPAL 95 and Alexander, 1998 private communication) for computing stellar interiors from the contracting Hayashi track to the white dwarf (or carbon ignition) phase. Photospheric modeling has reached high accuracy with theoretical models resembling very closely observed spectra from individual stars (\\pcite{Jimenez+98}). Furthermore, the new Hipparcos data allow for a much better calibration of the isochrones at the main-sequence (see \\pcite{Jimenez_Hipp+97} for more details) in conjunction with the classical calibration to the Sun. This then allows to construct self-consistent stellar population models that are properly calibrated to the Sun and to the Hipparcos data and therefore are more reliable than previous models. Also new stellar yields are now available (see section 4.1) and chemical evolution can be done more accurately. It is therefore, imperative and timely to explore the limits and predictions that models of stellar evolution make in light of the new data, such as chemical contents at early epochs, the evolution of star formation and luminosity density with time, as well as the spatial and other properties of the near-IR CIB and optical diffuse backgrounds. The outline of this paper is as follows: in Section 2 we discuss the cosmological context for the calculations we present in this paper. Section 3 discusses the stellar populations models we use for galaxy evolution. Modeling evolution of different galaxy types which at the same time is normalized to the data on modern galaxies and the redshift of galaxy formation is presented in Section 4. In Section 5 we compare our results to the data: we reproduce the evolution of chemical abundances at high $z$, the data on the luminosity density in $UV, B$ and $J$ bands from \\scite{Lilly+96}; the available data on deep galaxy counts in the various bands; and in in Section 5.4 compute properties of the near-IR CIB for the models that successfully reproduce the above data. Our conclusions are summarized in Section 6. ", + "conclusions": "In this paper we considered constraints from and predictions for the early Universe that follow from the evolutionary models of stellar populations in forming galaxies. We constructed accurate synthetic models for the evolution of stellar populations which are assumed to follow the Schmidt law for star formation. In modeling galaxy evolution we further account for chemical evolution. Early type galaxies have been normalized to the fundamental plane relations assumed to reflect variations in the mean metallicity along the luminosity sequence. Galaxy mixes are adopted from the CfA catalog. We assumed the Salpeter IMF for stars in early type galaxies and bulges of disk galaxies and the Scalo IMF for the disk stellar material. The galaxy numbers were normalized to the present-day galaxy luminosity function measurements in both $B$ and $K$ bands; both give consistent results. This allowed us to compute parameters that characterize the evolution of stellar populations in the early Universe to be compared with the available observational data. The computations were made for various cosmological density parameters $\\Omega, \\Omega_\\Lambda$ and the Hubble constant. Our main conclusions can be summarized as follows: 1) We computed the evolution of the mean cosmic metallicity with $z$. All models and cosmological parameters provide good fits to the current data. The current data have substantial uncertainties; after these get reduced one can hope to be able to further discriminate between the various models and cosmologies. 2) The evolution of the mean luminosity density in the $UV, B,J$ bands is well reproduced in our models, but models with zero cosmological constant are preferred. Our models reproduce well the recent HST data on the evolution of the star formation rate with redshift. 3) Our models give good fits to the available data on the deep galaxy counts in all bands where the observations are available, $B,R,I,K$, for low $\\Omega$ models (both flat and open Universe). The models thus do not require additional galaxy populations at intermediate and high redshifts in order to explain simultaneously $B$ and $K$ counts. We further fit well the redshift distribution of $B$ and $K$ counts in the magnitude range where such data are available. Our models give good fits to the color-magnitude data distribution, $I-K$ vs $K$, and do not produce extremely red colours even if early type galaxies form at very high redshifts. 4) We compute in detail the mean CIB flux produced by this evolution and the power spectrum of the CIB angular distribution. Our predictions are still below the current observational limits, but not by a large factor. This makes us optimistic that both the CIB and its angular structure can be measured in the upcoming years. We also computed the mean flux in the background light from this evolution in the visible bands, $B$ and $R$, and find that the recent positive measurements of the background at these bands require high redshift of galaxy formation, $z_F \\geq 5$. AK acknowledges support from NASA Long Term Space Astrophysics grant. \\newpage" + }, + "9802/astro-ph9802101_arXiv.txt": { + "abstract": "We present ROSAT HRI observations of six quasars with redshifts $\\ga$3.4: \\object{PKS 0335-122}, \\object{S4 0620+389}, \\object{B1422+231}, \\object{PKS 2215+02}, \\object{Q 2239-386} and \\object{PC 2331+0216}. We include the observation of the radio-quiet quasar \\object{BG 57 9}, whose redshift has recently been revised to $z = 0.965$. \\object{S4 0620+389}, \\object{B1422+231}, \\object{PKS 2215+02} and \\object{BG 57 9} are detected in the ROSAT energy band (0.1-2.4 keV). $2\\sigma$ upper limits are given for the remaining objects. All X-ray sources are point like to the limit of the HRI point spread function. We report marginal evidence for X-ray variability in \\object{S4 0620+389} on timescales of several hours (rest frame). No significant X-ray variability is found in the gravitationally lensed quasar \\object{B1422+231} on rest frame timescales of hours and months. We present the spectral energy distributions for six of the high redshift quasars and compare them to the mean distribution of low redshift quasars. ", + "introduction": "The study of quasars at high redshifts is motivated by a number of interesting, yet unsolved questions: How (and when) have quasars formed? How do their spectral energy distributions change with time? What causes the differences between radio-loud and radio-quiet quasars and do these differences persist to high redshifts? The X-ray emission of quasars is particularly interesting, since it constitutes a large fraction of the bolometric luminosity of quasars (e.g. Elvis et al. \\cite{elvis}). Further, the X-rays are mostly produced in the innermost regions of the AGN and can therefore provide clues to our understanding of the 'central engine' and quasar evolution. In the past, quasars with $z > 3$ have mostly been studied in the optical and the radio band. Only the advent of missions like ROSAT (Tr\\\"umper \\cite{truemper}) and ASCA (Tanaka et al. \\cite{tanaka}), which combine high sensitivity with sufficient spatial and spectral resolution, has enabled the investigation of the X-ray properties of these high redshift objects. However, the number of quasars with $z >3$, which are detected in X-rays, is still low (e.g. Cappi et al. 1997). In order to draw statistically significant conclusions on the emission properties of high redshift quasars and possible differences to their local counterparts, it is clearly desirable to increase the number of sources with X-ray data. We therefore observed with the ROSAT High Resolution Imager (HRI) four radio-loud (\\object{PKS 0335-122}, \\object{S4 0620+389}, \\object{B1422+231}, \\object{PKS 2215+02}) and two radio-quiet quasars (\\object{Q 2239-386}, \\object{PC 2331+0216}) at $z > 3.4$. The redshift of the radio-quiet quasar \\object{BG 57 9}, originally cataloged with $z = 3.74$ (Hewitt \\& Burbidge \\cite{hewitt87}), was recently revised to $z = 0.965$ (Borra et al. \\cite{borra}). Nevertheless, we present the X-ray observation in this paper. The basic properties of the observed quasars are summarized in Table \\ref{prop}. All quasars were previously not detected in X-rays. Due to its high spatial resolution and the achievable positional accuracy the HRI is very well suited to unambiguously identify the X-ray emission of high redshift quasars. ", + "conclusions": "We presented ROSAT HRI observations of six quasars with redshifts $z\\ga 3.4$ and one with $z = 0.965$. Four objects (\\object{S4 0620+389}, \\object{B1422+231}, \\object{PKS 2215+02}, \\object{BG 57 9}) are detected with X-ray luminosities ranging from $2.7 \\times 10^{44}$ erg~s$^{-1}$ to $5.1\\times 10^{46}$ erg~s$^{-1}$. 2$\\sigma$ upper limits are given for the remaining three objects (\\object{PKS 0335-122}, \\object{Q 2239-386}, \\object{PC 2331+0216}). All X-ray sources are point like to the limit of the ROSAT HRI point spread function ($\\approx$5\\arcsec, which corresponds to about 34 kpc at $z = 3.5$). We find marginal evidence for X-ray variability by a factor of two within the observation of \\object{S4 0620+389}, i.e. on a timescale of several hours in the rest frame of the source. The gravitationally lensed system \\object{B1422+231} shows no indication for X-ray variability during the individual observations as well as between the two HRI observations, which are separated by about one year. The spectral energy distributions of six of the high redshift quasars are presented. Although they are roughly consistent with the mean SED of low redshift quasars (Elvis et al. \\cite{elvis}), pronounced differences between individual objects appear. In particular \\object{PKS 2215+02} is almost two orders of magnitudes fainter in the optical compared to \\object{B1422+231}, while the radio and the X-ray luminosities differ by only a factor of two. However, the optical continuum slope measured by Francis et al. (in prep.) does not indicate significant extinction by dust." + }, + "9802/astro-ph9802193_arXiv.txt": { + "abstract": "We briefly discuss the theoretical implications of recent detections of gamma-ray bursts (GRBs) by BSAX. Relativistic shock wave theories of fireball expansion are challenged by the wealth of X-ray, optical and radio data obtained after the discovery of the first X-ray GRB afterglow. BSAX data contribute to address several issues concerning the initial and afterglow GRB emission. The observations also raise many questions that are still unsolved. The synchrotron shock model is in very good agreement with time-resolved broad-band spectra (2--500~keV) for the majority of GRBs detected by BSAX. ", + "introduction": "The discovery of X-ray afterglows by BSAX revolutionized the field of GRB research \\ci{costa,frontera,piro1}. The discovery cought us by surprise. Certainly, there was no theoretical prediction of hard \\ci{frontera,yoshida2} high-energy emission lasting hours-days after the main events. `Pre-BSAX models' of GRB shock waves predicted a much faster decay of the hard component produced by the initial impulsive particle acceleration (e.g., \\ci{mesz4}). Before the BSAX discovery, the only hope to detect GRB counterparts was believed to be searching for rapid UV/optical transients lasting a few minutes/hours or possibly delayed radio flares (e.g.,~\\ci{rhoads}). The ingenuity of the BSAX team was remarkable in ignoring theoretical models and carrying out the fastest ever slews of an X-ray satellite to GRB error boxes. BSAX discovery shattered old beliefs, and opened a new way of confronting difficult problems of GRB physics. A more complete picture of the GRB phenomenon is now emerging with all its complexities and puzzles. All theoretical models are challenged by the wealth of X-ray, ------------------------------------------------------\\\\ \\noindent {\\small ($\\star$) Paper presented at the Symposium {\\it The Active X-Ray Sky: Results from Beppo-SAX and Rossi-XTE}, Rome (Italy), 21-24 October 1997, Accademia Nazionale dei Lincei. To be published in Nuclear Physics B Proceedings Supplement, eds. L. Scarsi, H. Bradt, P. Giommi \\& F. Fiore. } \\noindent optical and radio data, and many problems remain unsolved at the moment. We briefly discuss\\\\ here some of the open issues. BSAX showed for the first time that a substantial fraction of GRB energy is dissipated in the X-ray range at late times (hours, days, sometimes weeks) after the main impulsive events. {\\it Are classical GRBs the tip of the iceberg of more complex radiation processes lasting much longer than expected from the cooling timescales of initial pulses ?} The `delayed' gamma-ray emission discovered by EGRET in the case of GRB~940217 \\ci{hurley} clearly shows how GRB `durations' determined by BATSE in the 50--300~keV range can be misleading. The delayed gamma-ray emission of GRB~940217 can now naturally be interpreted as a manifestation of GRB afterglows. Particles associated with this type of bursts can remain accelerated for a timescale $\\sim 10^2-10^3$ times longer than the decay timescales of initial GRB pulses. The fluence of the delayed gamma-ray emission of GRB~940217 is larger than $ 10\\%$ of that detected for the main event \\ci{hurley}. Also in the case of X-ray afterglows detected by BSAX, at least a fraction larger than $ 10\\%$ of the total fluence is emitted at late times (e.g., \\ci{costa,piro2}). The total energy inferred from the X-ray afterglow \\ci{piro2} and possible optical counterpart \\ci{pian} of GRB~970508 ($\\goe 10^{52}$~ergs) exceeds the most optimistic estimates of coalescing neutron star models (e.g., \\ci{paczinski0}). {\\it If GRB sources are at extragalactic distances, do they imply a new kind of explosive phenomenon ?} (see also the speculations of ref. \\ci{paczinskinew} to be compared with those of ref. \\ci{paczinski0}). Renormalizing the logN--logP distribution of GRBs (see Fig.~1) in terms of redshift for an assumed luminosity function can be done for GRB~970508 associated with the (lower limit) optical transient redshift $z=0.83$ \\ci{metzger1}. Unless GRB~970508 is anomalous, it can be shown that the standard candle assumption cannot satisfactorily describe the GRB brightness distribution, contrary to the conclusions of previous studies (e.g., \\ci{bloom}). Cosmological models of GRBs have to be formulated now without the appeal of a natural energy scale provided by the neutron star coalescence model. {\\it What is the origin of the spread in luminosity (and spectral characteristics as shown in ref. \\ci{tavani-logn}) of GRB sources ?} BSAX discoveries opened the way for rapid follow-up observations in the optical and radio bands. At present, four out of eight GRBs detected by BSAX (GRB~970228, 970402, 970508, 971214) unambiguously show the existence of fading X-ray sources within the WFC error boxes ($10-30 \\rm \\, arcmin^2$) pointed $\\sim 8 $~hours after the events. Of the remaining four GRBs, two error boxes could not be rapidly pointed (GRB~960720 and 980109), and the others pointed within $\\sim 14-16 $~hours show faint sources with ambiguous associations (GRB~970111 and 971227). Another GRB error box pointed by ASCA within 1 day after the event (GRB~970828) also shows an X-ray fading afterglow source \\ci{murakami,yoshida2}. Usually the X-ray flux decay is well represented by a power-law of the type $t^{-\\alpha}$. {\\it Why are the time exponents $\\alpha$'s of X-ray afterglows different from burst to burst~? Is there any correlation between the afterglow strength and the initial GRB peak intensity or spectrum ?} At present, three optical transients were identified within the WFC error boxes of GRB~970228 \\ci{vanparadijs}, GRB~970508 \\ci{bond} and GRB~971214 \\ci{halpern}. All are within the error boxes of the fading X-ray sources, reinforcing their associations with their respective GRBs. However, they all show different characteristics. The delayed ($\\sim 2 $~days) optical transient (OT) associated with GRB~970508 resulted in the spectral identification of absorption lines of an object at $z \\simeq 0.83$ (either the host or a foreground galaxy \\ci{metzger1}). After the delayed rise, the optical lightcurve follows a power-law decay of index $\\sim 1.17\\pm 0.04$ for several weeks (e.g., \\ci{pian}). The nature of the OT associated with GRB~970228 is currently controversial, with a nebulosity near a fading pointlike OT \\ci{vanparadijs,sahu,caraveo,fruchter}. The OT was still detectable by HST $\\sim 6$~months after the event near $V\\simeq 28$ \\ci{fruchter}. Its inferred optical lightcurve shows clear deviations from a power-law decay of index $\\sim 1.14\\pm 0.05$ \\ci{fruchter}. {\\bo The OT associated with GRB~971214 shows an initial power-law decay of index $\\sim 1.4\\pm 0.2$ \\ci{halpern}, and a possible flattening near $R = 25.6$ about $10$~days after the event \\ci{kulkarni}. {\\it What is the nature of the faint nebulosities associated with GRB~970228 and GRB~971214 ? If these are distant star-forming galaxies, why would GRBs be preferentially hosted near the cores of these galaxies rather than farther away as coalescing neutron star models predict ? If the comoving GRB formation rate follows that of star forming galaxies (strongly peaked near $z\\simeq 1$ \\ci{madau}), why are only the bright, harder and longer GRBs showing the strongest deviation from the Euclidean brightness distribution \\ci{tavani-logn} ?} Significant luminosity and spectral cosmological evolution of GRBs at $z\\goe 1$ may be required \\ci{tavani-logn}. \\begin{figure*} % \\vspace*{-3.cm} \\centerline{\\psfig{file=cumul.sax.aggiornato.ps,height=13.cm,width=13.cm,clip=}} \\vspace*{-.5cm} \\caption[image]{ GRBs detected by BSAX marked on the cumulative brightness distribution of the 4th BATSE catalogue \\ci{paciesas}. Peak intensities in the BSAX-GRBM detector have been rescaled to fit BATSE's energy range (50-300~keV) assuming a power-law photon index of 2. BSAX WFCs are clearly capable of detecting faint GRBs near the detection threshold for BATSE. Three GRBs associated with optical transients (OT) are marked with their respective peak magnitudes (GRB~970228 \\ci{vanparadijs}, GRB~970508 \\ci{castro}, GRB~971214 \\ci{halpern}). No OT was discovered in other GRB error boxes. The only GRB associated with a (possibly scintillating) radio source is GRB~970508 \\ci{frail} of low gamma-ray flux \\ci{piro2}. Radio searches in other GRB error boxes were rapidly performed with null results \\ci{frail3}.} \\label{fig1} \\end{figure*} About half of the GRBs promptly studied by optical searches do not show any OT below $R\\sim 22-23$ (GRB~970111, 970402, 970828, 970828, 971227, 980109). {\\it Is the lack of OTs in more than half of GRBs due to absorption within the host~? Is the lack of detectable absorption in many of the X-ray afterglow spectra obtained by SAX consistent with optical absorption [taking into account redshifting by $(1+z)^3$ of the X-ray cutoff energy]~?} GRB~970508, the event associated with an extragalactic OT, is `anomalous' in many ways when compared with other GRBs. {\\bf (i)} It is the only GRB with a non-monotonic X-ray afterglow decay $\\sim 1-3$ days after the event \\ci{piro2}. Four BSAX TOO observations were necessary to finally observe its decay in the 2--10~keV band (but apparently not in the softer LECS band \\ci{piro2}). {\\bf (ii)} Its associated OT was detected to rise $\\sim 2$~days after the event after a stable plateau state \\ci{bond}, contrary to the other two GRBs showing monotonically decreasing OTs (GRB~970228 \\ci{vanparadijs}, and GRB~971214 \\ci{halpern}). {\\bf (iii)} GRB~970508 is also the only GRB so far associated with a (scintillating~?) radio source \\ci{frail}. {\\it Are the peculiar properties of GRB~970508 caused by the `environment' surrounding the GRB source, or can they be attributed to a background AGN ? Why is only the very weak GRB~970508 associated with an apparently persistent \\ci{frail2} radio source~?} If delayed radio emission is common in bright GRBs produced by relativistic fireballs (e.g., \\ci{rhoads}), more radio detections would have been expected (see~Fig.~\\ref{fig1}). Finding satisfactory explanations to all these questions will be challenging for any theory of GRBs. The issues facing cosmological models \\ci{katzpiran,tavani,vietri,waxman,wijers} in the `post--BSAX era' of GRB research include: an energy crisis, burst number density evolution vs. star forming galaxies, luminosity and spectral evolution, absorption and reprocessing properties of GRB environments in distant galaxies, diversity of X-ray and optical decays, persistent optical emission $\\sim 6$~months after GRB~970228, and lack of radio emission for the majority of GRBs. Other models should be considered, and a superposition of GRB populations of different origins may still be viable. More detections of optical/radio transients in GRB error boxes are definitely needed. ", + "conclusions": "" + }, + "9802/astro-ph9802250_arXiv.txt": { + "abstract": "We have performed a series of simulations of clusters of galaxies on the basis of the smoothed particle hydrodynamics technique in a spatially-flat cold dark matter universe with $\\Omega=0.3$, $\\lambda=0.7$, and $H_0=70$km/s/Mpc as one of the most successful representative cosmological scenarios. In particular, we focus on the Sunyaev--Zel'dovich effect in submm and mm bands, and estimate the reliability of the estimates of the global Hubble constant $H_0$ and the peculiar velocity of clusters $v_r$. Our simulations indicate that fractional uncertainties of the estimates of $H_0$ amount to $\\sim 20$\\% mainly due to the departure from the isothermal and spherical gas density distribution. We find a systematic underestimate bias of $H_0$ by $\\sim 20$\\% for clusters $z\\approx 1$, but not at $z\\approx 0$. The gas temperature drop in the central regions of our simulated clusters leads to the underestimate bias of $v_r$ by $\\sim 5$\\% at $z\\approx 0$ and by $\\sim 15$\\% at $z\\approx 1$ in addition to the statistical errors of the comparable amount due to the non-spherical gas profile. ", + "introduction": "Clusters of galaxies have been extensively observed in radio, optical and X--ray bands. Furthermore recent and future observational facilities in mm and submm bands, such as the SCUBA (Submillimeter Common-User Bolometer Array), the Japanese LMSA (Large Millimeter and Submillimeter Array) project and the European PLANCK mission are expected to open the submm window to observe clusters of galaxies via the Sunyaev--Zel'dovich (SZ) effect (Sunyaev, Zel'dovich 1972) in addition to the Rayleigh-Jeans region of the spectrum of the cosmic microwave background (CMB) where the SZ temperature decrement is reported for about a dozen of clusters (e.g., Rephaeli 1995; Kobayashi, Sasaki, Suto 1996). Since the intensity of the SZ effect does not suffer from the $(1+z)^{-4}$ diminishing factor unlike X--ray surface brightness, observations in mm and submm bands are much more advantageous for clusters, especially at high $z$, than those in optical and X-ray bands (Barbosa et al. 1996; Silverberg et al. 1997; Kitayama, Sasaki, Suto 1998). As extensively discussed in previous literatures (Silk, White 1978; Sunyaev, Zel'dovich 1980; Birkinshaw, Hughes, Arnaud 1991; Rephaeli, Lahav 1991; Inagaki, Suginohara, Suto 1995; Kobayashi, Sasaki, Suto 1996; Holzapfel et al. 1997; Kitayama, Sasaki, Suto 1998), one can then combine these multi-band observations of high--$z$ clusters to determine the cosmological parameters and the peculiar velocity of clusters. These procedures, however, usually assume that the gas of clusters is isothermal and spherical, while all the observed clusters do exhibit a departure from the assumption to some extent. The departure would hamper the reliable estimates of, for instance, the Hubble constant $H_0$ and the peculiar velocity $v_r$. To address the question quantitatively, we carry out a series of numerical simulations of clusters. We extract simulated clusters both at $z\\approx 0.0$ and $z\\approx 1.0$, and perform the ``simulated'' observations in X--ray, mm and submm bands. Finally we combine the multi-band information to evaluate the statistical and possible systematic errors of the estimates of $H_0$ and $v_r$. Our present work extends the previous studies of this methodology (Inagaki, Suginohara, Suto 1995; Roettiger, Stone, Mushotzky 1997) for $H_0$ and examines uncertainties of the peculiar velocity field as well, paying attention to the projection effect and the evolution of clusters. ", + "conclusions": "" + }, + "9802/astro-ph9802299_arXiv.txt": { + "abstract": "We present a spectrum of the cool ($T_{\\rm eff} = 900$ K) brown dwarf Gliese 229B. This spectrum, with a relatively high signal-to-noise ratio per spectral resolution element ($\\gtrsim 30$), spans the wavelength range from 0.837 $\\mu$m to 5.0 $\\mu$m. We identify a total of four different major methane absorption features, including the fundamental band at 3.3 $\\mu$m, at least four steam bands, and two neutral cesium features. We confirm the recent detection of carbon monoxide (CO) in excess of what is predicted by thermochemical equilibrium calculations. Carbon is primarily involved in a chemical balance between methane and CO at the temperatures and pressures present in the outer parts of a brown dwarf. At lower temperatures, the balance favors methane, while in the deeper, hotter regions, the reaction reverses to convert methane into CO. The presence of CO in the observable part of the atmosphere is therefore a sensitive indicator of vertical flows. The high signal-to-noise ratio in the 1 $\\mu$m to 2.5 $\\mu$m region permits us to place constraints on the quantity of dust in the atmosphere of the brown dwarf. We are unable to reconcile the observed spectrum with synthetic spectra that include the presences of dust. The presence of CO but lack of dust may be a clue to the location of the boundaries of the outer convective region of the atmosphere: The lack of dust may mean that it is not being conveyed into the photosphere by convection, or that it exists in patchy clouds. If the dust is not in clouds, but rather sits below the outer convective region, we estimate that the boundary between outer convective and inner radiative layers is between 1250 K and 1600 K, in agreement with recent models. ", + "introduction": "Gliese 229B is the first sub-stellar object outside the solar system with an effective temperature, T$_{\\rm eff}$, well below 1800 K, the minimum T$_{\\rm eff}$ for stars, found using direct imaging techniques (Nakajima et al. \\cite{nak95}). Other sub-stellar companions have been inferred to exist through indirect techniques (Wolszczan \\&\\ Frail \\cite{wf92}; Mayor \\&\\ Queloz \\cite{mq95}; Butler et al. \\cite{but97}). The astrophysical nature of all these objects can only be understood in detail through spectroscopic studies. Unfortunately current technology precludes the spectroscopic study of any of these objects, except Gliese 229B. Interestingly, the spectrum of a gaseous, sub-stellar object is primarily determined by its effective temperature, and secondarily by its surface gravity and composition (Burrows et al. \\cite{bur97}; Burrows et al. \\cite{bur96}). Thus, the spectrum of Gliese 229B is expected to resemble the spectra of some of the less massive, but similarly hot, giant planets found in the radial velocity studies. Here we present the spectrum of Gliese 229B with high signal to noise ratio from 0.837 $\\mu$m out to 5 $\\mu$m. Gliese 229B emits approximately 65\\% of its emergent flux in this wavelength region. An effort like this can perhaps be seen as a precursor to extrasolar planet studies or an extension of the work of hundreds of planetary scientists who have studied and modelled the atmospheres of the planets orbitting the Sun. In addition, these observations also give extremely interesting insights into the atmospheric physics of brown dwarfs, including the atmospheric dynamics. Our now detailed set of observations can be compared with recent models. ", + "conclusions": "" + }, + "9802/astro-ph9802120_arXiv.txt": { + "abstract": "The C~I, Na~I~D, and H$\\alpha$ lines of the post-AGB binary HR~4049 have been studied. Na~I~D variability results from a photospheric absorption component ([Na/H]=$-1.6\\pm0.2$) which follows the velocity of the primary and a stationary, non-photospheric component. An emission component is attributed to the circumbinary disc, and an absorption component to mass-loss from the system with a velocity of $5.3\\pm0.5$~\\kms. The H$\\alpha$ profile varies with the orbital period. The two strong shell type emission peaks are identified as from one single broad emission feature with an absorption centered around $-7.5$~\\kms. The intensity variations are largely attributed to a differential amount of reddening towards the H$\\alpha$ emitting region and the stellar continuum. The radial velocities suggest that the H$\\alpha$ emission moves in phase with the primary, but with a slightly lower velocity amplitude. From this we infer that the H$\\alpha$ emission comes from outside the orbit of the primary, but still gravitational bound to the primary. H$\\alpha$ also shows a weak emission feature at $-21.3\\pm3.5$~\\kms, which originates from the circumbinary disc and a weak absorption feature at $-7.5\\pm1.6$~\\kms~ due to absorption by the circumbinary disc. We propose two competing models that could account for the observed velocity and intensity variations of the H$\\alpha$ profile. Model~I: light from the primary reflects on a localized spot near the inner radius of the circumbinary disc which is closest to the primary. Model~II: H$\\alpha$ emission originates in the outer layers of the extended atmosphere of the primary due to activity. These activities are locked to the position of the primary in its orbit. We discuss the similarities of variability and shape of the H$\\alpha$ emission of HR~4049 with those of early type T-Tauri stars (e.g SU~Aur). ", + "introduction": "With [Fe/H]~$\\approx -4.8$, HR~4049 is the prototype of a group of high-latitude supergiant binary stars with an extremely metal-depleted photosphere (van Winckel et al. \\cite{winckelwaelkens} and references therein). Their high luminosity, position above the Galactic plane and in most cases observationally confirmed C-rich circumstellar dust, suggest that they are in a post-AGB phase of evolution. The observed photospheric abundance patterns (low abundances of chemical elements with high dust condensation temperature and high abundance of elements with low dust condensation temperature) can best be explained by a model in which circumstellar gas, devoid of refractory elements, is accreted on the star, coating it with a chemically peculiar layer (Venn \\& Lambert \\cite{vennlambert}; van Winckel et al. \\cite{winckelmathis}). For the extremely metal-depleted post-AGB stars, the gas-dust separation occurs in a circumbinary disk (Waters et al. \\cite{watersetal}) but the question of the efficiency of the process gained momentum after the findings of Giridhar et al. (\\cite{giridharetal}), Gonzalez et al. (\\cite{gonzalezetala} and \\cite{gonzalezetalb}), and Gonzalez \\& Wallerstein (\\cite{gonzalezwallerstein}) that also the photosphere of Galactic RV~Tauri stars and the binary type~II Cepheid ST~Pup show the same depletion pattern. HR~4049 is a binary system with an orbital period of 430 days (Waelkens et al. \\cite{waelkenslamers}). All other members of the group of metal-depleted post-AGB binaries have orbital periods of the order of one to several years (van Winckel et al. \\cite{winckelwaelkens}). With an estimated stellar radius of 47~R$_{\\odot}$, the primary of HR~4049 fits nicely within the binary system. However, single star evolution predicts a stellar radius of about 250~R$_{\\odot}$ at the tip of the AGB (Boothroyd \\& Sackmann \\cite{boothroydsackmann}). Since the minimum separation between the two stars is 190~R$_{\\odot}$ this would lead to a phase of common envelope evolution (case C Roche lobe overflow). Evidently, the system survived this phase, or did not experience it. Either way, we expect that the stellar masses of the two bodies might have been altered. HR~4049 shows photometric variations which are attributed to a phase-dependent circumsystem reddening (Waelkens et al. \\cite{waelkenslamers}) from a circumbinary disc (CBD). The inclination of the system is not known, but both edge-on and face-on are excluded since the circumsystem reddening is phase-dependent. Supporting evidence for a CBD comes from the presence of large dust grains ($> 0.8$~$\\mu$m) as inferred from the polarimetric measurements by Joshi et al. (\\cite{joshietal}), and the presence of relatively hot dust close to the star (Lamers et al. \\cite{lamersetal}). In order to understand the environment of HR~4049, we have made a detailed study of the variability of the C~I, Na~I~D and H$\\alpha$ lines of HR~4049. In Sect.~2 we present the observations and in Sect.~3 re-derive the orbital parameters. The variability of the Na~D lines is described in Sect.~4 and of H$\\alpha$ in Sect.~5. In Sect.~6 we propose two competing models to account for the H$\\alpha$ variability. A discussion is presented in Sect.~7. ", + "conclusions": "The H$\\alpha$ and Na~I~D line profiles of HR~4049 are clearly varying with the orbital period. The C~I line profiles are asymmetric but show no variations with orbital phase. The asymmetry is likely due to granulation. The changes in the observed line profiles of Na~I~D are due to the Doppler shift of the photospheric component relative to the circumstellar and interstellar emission and absorption spectrum. There are several interstellar absorption components (B, C$_{1}$, C$_{2}$, and E$_{2}$), a broad absorption from the circumbinary disc (A$_{2}$), emission from the approaching part of the CBD (A$_{3}$), and a weak mass-loss component (A$_{1}$) at a velocity of only $5.3\\pm0.5$~\\kms. For H$\\alpha$ we identify photospheric component, only detectable as extended absorption wings (Fig.~\\ref{fig_halphamodel}). The CBD absorbs (R$_{\\rm min}$) and emits (C$_{\\rm max}$) in H$\\alpha$. The absorption from the CBD is 7~\\kms~ blue-shifted from the systemic velocity. If this shift is real, it could suggest that we see a photospheric spectrum reflected by circumstellar material. In such a scenario the systemic velocity derived from the photospheric spectrum represent the velocity component in the direction of the mirror and not towards the observer. Furthermore, the mirror could have a small velocity. The critical question here is what is the systemic velocity from CO radio observations. Unfortunately no such data is available. There is a weak emission feature at $-21.3\\pm3.5$~\\kms~ relative to the systemic velocity, which originates from the circumbinary disc. The two strong shell type H$\\alpha$ emission peaks are from one single broad emission feature with a central absorption feature centered around $-7.5$~\\kms~ relative to the center of the emission. The intensity variations are best explained as due to a differential amount of reddening towards the H$\\alpha$ emitting region and the stellar continuum, with a possible second order variation such that H$\\alpha$ emission is stronger near periastron than apastron. The radial velocities suggest that the H$\\alpha$ emission moves in phase with the primary, but with a slightly lower velocity amplitude. We proposed two competing models. In both models the H$\\alpha$ profile is interpreted as a single broad emission feature with a central absorption, and the relative intensity variations are attributed to a different amount of reddening toward the H$\\alpha$ emitting volume and the stellar continuum. In model~I, H$\\alpha$ emission is from reflected starlight on a localized spot in the circumbinary disc. In model~II, the H$\\alpha$ emission originated from the extended atmosphere of the primary. \\begin{figure} % \\centerline{\\hbox{\\psfig{figure=art14halphamodel.ps,width=\\columnwidth,angle=90}}} \\caption{Composite of the H$\\alpha$ profile. From top to bottom (with increasing distance from the star): The photospheric profile of the primary, the double peaked emission (DPE), emission and absorption from the CBD, and the composite profile. This combined profile depends on the phase since the photospheric spectrum is experiencing a phase-dependent reddening.} \\label{fig_halphamodel} \\end{figure} There is considerable similarity between the HR~4049's H$\\alpha$ profile and that of SU~Aur, an early-type T Tauri star (ETTS, Herbst et al. \\cite{herbstetal}, other members are CO Ori, RY Tau, and RY Lup) with an accretion disc viewed almost edge on - see H$\\alpha$ profiles published by Giampapa et al. (\\cite{giampapaetal}), Johns \\& Basri (\\cite{johnsbasri}), and Fig.~3 of Petrov et al. (\\cite{petrovetal}). In both stars, the red emission is more pronounced than the blue emission with a central absorption near the systemic velocity. It is this similarity that drew our attention to the models proposed to account for the H$\\alpha$ and other Balmer line profiles. Although the H$\\alpha$ profiles of HR~4049 and SU Tau are similar in form, they are different in scale: the emission for SU Tau is spread over a velocity range that is about eight times that seen here for HR~4049. This difference is likely traceable to the higher surface gravity of SU Tau and the related differences in escape velocity ($v \\approx 70$~\\kms for HR~4049 but $v \\approx 511$~\\kms~ for SU Tau) or in Keplerian velocity at a fixed distance from the star. Giampapa et al. (\\cite{giampapaetal}) proposed a two component wind model for SU~Aur: an optically thin wind moving at high velocity outside a slow, optically thick wind. Johns \\& Basri (\\cite{johnsbasri}) proposed an alternative model with a slow moving (or stationary) region close to the star with a very high turbulence velocity, and a wind at a terminal velocity of about 50 \\kms. In this model, the wind originates in the disc at the place where the disc's Keplerian angular velocity is equal to the star's angular rotation rate (typically at a distance of $3~R_{A}$). The central absorption in SU~Aur is from the stellar wind which leaves the disc at 150 \\kms~ and decelerate due to the gravitational force of the stars and absorbs at a terminal velocity of 50~\\kms. Interestingly, the photometry varies non-periodically and it has been suggested that this is due to variable circumstellar extinction. From a Fourier analysis of line profiles, Bouvier et al. (\\cite{bouvier}) claim that SU~Aur is a double lined spectroscopic binary, but the binary nature of SU~Aur has not been confirmed using speckle imaging (Ghez et al. \\cite{ghez}). Based on presence of a disc and the similar shape of the H$\\alpha$ emission, it is tempting to believe that the emitting region of SU~Aur might resemble that of HR~4049. It is interesting to speculate in which respect our proposed models for H$\\alpha$ emission could help to understand the gas-dust separation process. First of all, it seems that all metal-depleted post-AGB binaries have H$\\alpha$ emission (from very weak (BD~+39$^{\\circ}$4926) to extremely strong (HR~4049 and HD~44179)) that is periodic with the orbital period (this is confirmed for HR~4049 and HD~213985) It is not evident that the H$\\alpha$ emission is related to the metal-depleted nature of the star. BD~+39$^{\\circ}$4926 which is extremely metal depleted has very little H$\\alpha$ emission, but also post-AGB stars which are not in a binary system show H$\\alpha$ emission which is variable (HD~56126 and HD~133656). Therefore we argue that the variability of the H$\\alpha$ emission cannot be linked directly to the metal-depleted nature of the post-AGB stars. \tInstead it seems that the H$\\alpha$ emission and the metal-depleted photosphere are both the result of the binary nature of these stars and the geometry of their circumstellar environment. Based on our analysis we suggest to proceed with a detailed study on the shape of the bisector with orbital phase. This could lead to a better understanding of the effect of granulation in post-AGB stars. Spectropolarimetric observations of H$\\alpha$ and other emission lines could distinguish between reflected and direct light. Finally, the selective reddening scenario can be tested with simultaneous high-resolution spectroscopy and photometry." + }, + "9802/astro-ph9802316_arXiv.txt": { + "abstract": "Broad-band (gamma to radio) variations of the flux density were observed in the first half of 1992 in the luminous high redshift ($z=2.172$) quasar S5 \\object{0836+710}. VLBI monitoring observations during 1993 -- 1996 performed at 86\\,GHz, 22\\,GHz, 15\\,GHz, and 8\\,GHz show the ejection of a new jet component, which most probably is directly related to a quasi simultaneous gamma--, X-ray, optical flaring activity which was observed in February 1992. During the period 1992 -- 1993 the flaring propagated through the radio spectrum. From several quasi-simultaneous radio spectra taken during this phase of activity, we determine the time evolution of the spectral turnover of the radio spectrum in the $S_{\\rm m}$-$\\nu_{\\rm m}$ diagram. The data indicate a correlation of the jet activity with the variability of the broad-band electromagnetic spectrum of the source. The observational findings are discussed in the framework of relativistic shock models. ", + "introduction": "} In early 1992, the ultraluminous S5 quasar \\object{0836+710} (\\object{4C\\,71.07}, $z=2.172$) underwent a prominent optical outburst (von Linde et al.\\ 1993), which later led to enhanced variability in the mm- and cm- radio bands (\\cite{mar94}). At the time of the optical flaring, the source also was in a bright state in the gamma-regime (EGRET: Fichtel et al.\\ 1994) and at soft X-rays (ROSAT: \\cite{bru94}). Motivated by the increased activity of the source, we started a high frequency VLBI monitoring program in 1993 in order to investigate the possible relation between broad-band changes in the flux density and structural variations in the jet on sub--milliarcsecond scales. The ejection of new VLBI jet components after major outbursts in the flux density has been observed also in a few other sources (e.\\ g.\\ component ejection after optical flaring: \\object{3C\\,345} Babadzhanyants and Belokon 1986, \\object{3C\\,273} \\cite{kri90a}; ejection after radio flares: \\object{BL\\,Lac} Mutel et al.\\ 1990, \\object{PKS\\,0420$-$014}, Wagner et al.\\ 1995a) and seems to be a quite common phenomenon in compact flat spectrum radio sources (blazars). With the new data obtained for \\object{0836+710}, it appears very likely that the relation between outbursts in the flux density observed at high frequencies and jet component ejection is not limited to the synchrotron regime (e.g. optical-to-radio), but most probably covers a much wider spectral range (gamma-to-radio), possibly the full electromagnetic spectrum. Similar correlations between gamma-activity and jet activity were recently found also in a number of other sources, e.g. in 3C~273 (Krichbaum et al.\\ 1996), 3C~279 (Wehrle et al.\\ 1994), 0528+134 and 3C454.3 (Krichbaum et al.\\ 1995), where the ejection of jet components could be related to gamma-flares, or enhanced levels of gamma-activity. However, in most of the published cases the frequency coverage is not as broad as presented here for \\object{0836+710}. The quasar \\object{0836+710} is a member of a complete flux density limited sample of 13 flat--spectrum radio sources (Eckart et al.\\ 1986 \\& 1987) compiled from the 6\\,cm S5 survey (K\\\"uhr et al.\\ 1981). The 13 sources of the sample are located in the northern polar cap ($\\delta \\geq 70\\deg$) and are regularly studied with VLBI at frequencies ranging from 0.3 to 22\\,GHz (cf. Witzel et al.\\ 1988 and references therein). The VLBI data for \\object{0836+710} now span a time range of more than 15 years (1983--1996). The high dynamic range VLBI images of \\object{0836+710} reveal a complex and wiggled one--sided core--jet structure (Eckart et al.\\ 1986, \\cite{eck87}). World-array VLBI images at 90\\,cm and 18\\,cm show the jet extending over more than 150\\,milliarc\\-seconds (mas) in direction to the outer arcsecond lobe seen with MERLIN and the VLA (\\cite{hum92}). Mainly from the VLBI monitoring at 6\\,cm (Krichbaum et al.\\ 1990b hereafter K90b, Otterbein 1996 hereafter O96), a very complex motion pattern is revealed, with jet components, moving at velocities ranging from subluminal ($\\beta_{\\rm app}=0.5; \\beta = v/c$) to superluminal ($\\beta_{\\rm app}=10$) (assuming $\\rm H_0 = 100\\, \\rm{km/s/Mpc}$ and $q_0 = 0.5$, throughout this paper). No systematic correlation between the speed of the jet and the separation of the components from the core seems to be present. The jet shows a noticeably structure, with lateral displacements of its ridge line (kinks) and oscillations of its transverse width (K90b). At cm-wavelengths, \\object{0836+710} is strongly polarized ($\\rm p \\simeq 9\\,\\%$ at 5\\,GHz). Polarization VLBI images show a highly polarized jet with the magnetic field roughly following the bent ridge line (\\cite{caw93}). This suggests that the jet of \\object{0836+710} is also highly magnetized. In this paper, we focus on the structural changes observed in the sub-milliarcsecond regions of the inner jet of \\object{0836+710} during and after the optical outburst of 1992. A more detailed discussion of the overall morphology and the spectral and kinematic properties of the milliarcsecond-jet will be given elsewhere. \\begin{table} {\\small \\caption{\\label{sum} Summary of experiments} \\begin{tabular}{lcll} date & freq. &observing mode, & array \\\\ & [GHz] &bandwidth [MHz] & \\\\ \\hline 1993, Apr. 7 & 86.2 &MK\\,III (A), 112 & global (4)$^1$ \\\\ 1993, Sep. 16 & 22.2 &MK\\,III (B), 56 & global (11) \\\\ 1995, Jan. 20 & 15.3 &VLBA 64-8-1, 32 & VLBA (10)$^2$ \\\\ 1995, Aug. 24 & 22.2 &VLBA 128-8-1,64 & VLBA (10)$^3$ \\\\ 1995, Aug. 24 & ~8.4 &VLBA 128-8-1,64 & VLBA (10)$^3$ \\\\ \\end{tabular} \\vspace{0.3cm} Notes: \\\\ $^1$: stations: Effelsberg, Pico Veleta, Onsala, Haystack\\\\ $^2$: snapshot--type observations \\\\ $^3$: 8/22\\,GHz simultaneous dual--frequency observations \\\\ } \\end{table} ", + "conclusions": "} \\subsection{The total spectrum and the VLBI jet \\label{sc:model}} The $S_{\\rm m}$--$\\nu_{\\rm m}$ dependence determined in the previous section can be used for studying the relation between the spectral and kinematic properties of the relativistic jet in \\object{0836+710}. Despite some evidence for emission from an inhomogeneous synchrotron component ($\\alpha_{\\rm thick} = 1.0$, see previous section), we postulate for the sake of simplicity, that the observed spectral variations are produced by a single, compact relativistic shock dominating the source radio emission, and associated with the core of the jet or an emitting region moving inside the jet within $\\lessim 0.1$\\,mas distance from the core (a distance comparable to the back--extrapolated separation of B3 at the last spectral epoch, 1992.9). We use the formalism developed by Marscher \\& Gear (1985), and account for possible variations of the Doppler factor in the emitting region (Marscher 1990). Based on the synchrotron spectral index $\\alpha=-0.7$ used for the spectral fitting, a power law electron energy distribution with $s=2.4\\,\\, (s=1-2\\alpha)$ is assumed. The magnetic field distribution along the jet axis $r$ is defined by an exponent $a$, so that $B \\propto r^{-a}$. The Doppler factor is $\\delta \\propto r^b$. The spectral evolution of the shock is then described by the variations of the turnover point: $S_{\\rm m}\\propto \\nu_{\\rm m}^{\\rho}$ and $\\nu_{\\rm m} \\propto r^{\\epsilon}$. The exponents $\\rho$ and $\\epsilon$ depend on the dominating type of the energy losses (Marscher 1990, Marscher et al.\\ 1991, Lobanov \\& Zensus 1998). Following Marscher (1990), we obtain the following expressions for $\\rho$ and $\\epsilon$ during the adiabatic--loss stage: \\begin{equation} \\label{eq:index_Sm} \\rho = -\\frac{(19-4s) - 3a(2s+3) + 3b(3s+7)}{2(2s+1)+3(a-b)(s+2)}\\,, \\end{equation} \\begin{equation} \\label{eq:index_num} \\epsilon = -\\frac{2(2s+1)+3(a-b)(s+2)}{3(s+4)}\\,. \\end{equation} We now consider two cases: a)~a jet with transversal ($a=1$) magnetic field, and b)~a jet with longitudinal ($a=2$) magnetic field. With the measured $\\rho=0.58$, changes of the Doppler factor and turnover frequency along the jet can be determined from (\\ref{eq:index_Sm}) and (\\ref{eq:index_num}). For $a=1$, we obtain $b=0.0$ and $\\epsilon=-1.3$. The distance travelled by the emitting region between 1992.07 and 1992.90 (the 1st and the 7th spectral epochs) should increase by a factor of $r_7/r_1\\approx 2.1$. To assess the jet kinematics, we use $\\beta_{\\rm app} = 10.8$ measured in Section \\ref{sc:B3} for the jet component B3 at small separations from the core. For now, we assume that the jet bulk Lorentz factor is $\\gamma_{\\rm j}= 10.9$, with the resulting Doppler factor $\\delta_{\\rm j} = 11.4$ and jet viewing angle $\\theta_{\\rm j}=5\\deg$. Using the time separation between the two aforementioned spectral epochs, $\\Delta t = 0.83$\\,yrs, one can then estimate the distance between the emitting region and the jet apex, $r_1 = \\beta \\gamma \\delta \\Delta t / (1+z) = 9.5\\pm0.8$\\,pc. The case with $a=2$ results in $b=0.4$, implying that the emitting region was accelerating. The corresponding turnover frequency evolution is described by $\\epsilon=-1.7$. The resulting Doppler factor must increase between $t_1$ and $t_7$ by $\\delta_7/\\delta_1 = (\\nu_{\\rm m\\,7}/\\nu_{\\rm m\\,1})^{b/\\epsilon} \\approx 1.3$, and the corresponding distance must become larger by a factor of $\\approx 1.8$. For a jet with variable Doppler factor, the distance travelled by the shock between two epochs $t_1$ and $t_2$ is (Lobanov \\& Zensus 1998): \\begin{equation} \\label{eq:travdistance} \\Delta r_{1,2} = (1+z)^{-1} \\int_{t_1}^{t_2} \\frac{\\beta(t) {\\rm d}t} {1 - \\beta(t) \\cos\\theta(t)}\\,. \\end{equation} For the Doppler factor $\\delta \\propto r^b$, equation (\\ref{eq:travdistance}) gives for the distance to the jet apex at the epoch $t_1$: \\begin{equation} \\label{eq:zerodistance} r_1 = \\left( \\frac{1+z}{\\delta_1 c \\Delta t} \\int_{1}^{r_u} \\frac{1}{\\sqrt{\\gamma^2(r)-1}} \\frac{{\\rm d}r}{r^b} \\right)^{1/(b-1)}\\,, \\end{equation} where $c$ is the light speed, $\\Delta t = t_2 - t_1$, and $r_u = (\\nu_{\\rm m\\,2}/\\nu_{\\rm m\\,1})^{1/\\epsilon}$. The exact form of $\\gamma(r)$ is essentially unknown. In many cases, it can be assumed to be constant (so that the Doppler factor variations are entirely due to a curved path followed by the emitting region). For a straight jet, the $\\gamma(r)$ can be derived from the Doppler factor variations. For \\object{0836+710}, both approaches yield similar result: $r_1 \\sim 30$\\,pc (here we postulate $\\delta_7 = \\delta_{\\rm j}$). One can argue that the location $r_1$ (the beginning of the adiabatic--loss stage) should correspond to the VLBI core, since the core emission is likely to be dominated by Compton and synchrotron energy losses (Unwin et al.\\ 1994, 1997), whereas the moving jet components are successfully modelled as adiabatically expanding spherical plasmons embedded in the jet (Zensus et al.\\ 1995). In view of this argument, the case of a non--accelerating jet with $r_1\\approx 10$\\,pc appears to be more likely to explain the observed spectral behavior in \\object{0836+710}. \\subsection{Travel times and radio--to--gamma correlations \\label{sc:travel}} If the jet core is located at $r_1$, one can calculate the travel time between the jet nozzle, $r_0$, and the core. For $b=0$ ($\\delta=$const), the calculation is trivial, and it gives (in the observer's frame) $\\Delta t_{0,1} = (1+z) r_1 \\gamma^{-1} \\delta^{-1} \\beta^{-1} = 0.74\\pm0.07$\\,yrs. For an accelerating jet, the travel time can be calculated from (\\ref{eq:travdistance}), and it depends on the (unknown) Doppler factor $\\delta_0$ at the jet nozzle. If the nozzle is formed by a pressure drop in a relativistic outflow, the resulting Lorentz factor is $\\gamma_0 \\approx 1.2$ (Marscher 1980). The corresponding Doppler factor is then $\\delta_0 \\approx 1.8$, for $\\theta_0 = \\theta_{\\rm j} = 5\\deg$. The resulting travel time is $\\Delta t_{0,1} = 1.6$\\,yrs. For larger $\\gamma_0$, the travel time increases and may become as large as 2.8\\,yrs, if $\\gamma_0 \\approx \\gamma_1$. On the basis of the arguments presented in the previous section, we take $\\Delta t_{0,1} = 0.74$\\,yrs as a better estimate of the travel time between the nozzle and the core. In section \\ref{sc:B3}, we have found $t_{\\rm ej}=1992.65$ for the ejection time of the superluminal feature B3. Together with the derived nozzle--to--core travel time, this implies that the plasma condensation responsible for the emission of B3 had travelled through the nozzle at $t_0 = 1991.91\\pm0.21$. The latter epoch is very close to the time when the optical flare was observed in the source. Although the coincidence may be only fortuitous, it is rather striking, and we would like to emphasize that such a situation is indeed possible. For instance, if a flare occurs when a dense plasma condensation travels through the jet sonic point, the optical depth in the radio bands can exceed unity, and only a high--energy flare (optical, X--rays, or gamma--rays) is observed. An increase of the radio emission and the associated ejection of a new VLBI feature are observed later, after the condensation has travelled outside the $\\tau\\geq 1$ region. The jet geometry and travel times above are determined for $ \\gamma_{\\rm j}\\approx\\gamma_{\\rm min}$. If we suppose that the gamma-to-optical flaring activity and the ejection of VLBI component B3 in 1992 had a common origin, the jet geometry can be uniquely determined. We assume that the maximum of the flare occurs at $t_{\\rm flare} = 1992.13$ and take $t_{\\rm ej} = 1992.65$ from the extrapolation of the component motion. In such a scenario, $\\Delta t_{0,1} = t_{\\rm ej} - t_{\\rm flare}$, and we find $\\gamma_{\\rm j} = 11.8$ and $\\theta_{\\rm j} = 3\\fdg2$ . The corresponding Doppler factor is $\\delta_{\\rm j} = 16.4$; the core--to--nozzle distance is $r_1 \\approx 15$\\,pc. We note that the derived values are consistent with overall kinematics of \\object{0836+710} (O96). We therefore view this model as a likely explanation for the kinematic properties, spectral behavior and flaring activity observed in \\object{0836+710} in 1992--1993. \\subsection{Properties of the inner jet \\label{sc:inner-jet}} The kinematic parameters determined above for the case of a non--accelerating ($b=0$) jet with transverse ($a=1$) magnetic field can be further used to evaluate the physical conditions in the emitting region responsible for the observed spectral changes in \\object{0836+710}. For this purpose, we associate the core--to--nozzle distance, $r_1 = 15$\\,pc with the location of the emitting region at the estimated epoch of origin of B3 ($t_{\\rm ej} = 1992.65$). Then, for each spectral epoch, we find the respective location, \\begin{equation} r = r_1 \\left(\\frac{\\nu_{\\rm m}}{\\nu_1}\\right)^{1/\\epsilon}\\, , \\label{eq:jet5} \\end{equation} and size of the emitting region (using equation (3) from Marscher 1987), \\begin{equation} d = d_1 \\left(\\frac{\\nu_1}{\\nu_{\\rm m}}\\right)^{(5\\epsilon + a+b )/4\\epsilon} \\left(\\frac{S_{\\rm m}}{S_1}\\right)^{1/2} \\, . \\label{eq:jet6} \\end{equation} $S_1$ and $\\nu_1$ are measured at $t_{\\rm ej}$, which gives $\\nu_1 = 17.9\\pm1.8$\\,GHz and $S_1 = 0.98\\pm0.11$\\,Jy. For $d_1$, we use the size of the VLBI core, $\\Omega_{\\rm core}\\approx 0.1$\\,mas, typically measured in the 22\\,GHz images of \\object{0836+710}. The corresponding linear size is $d_1 = 0.4$\\,pc. \\begin{figure}[t] \\mbox{\\psfig{figure=7484f9.ps,height=6.5cm,width=8.5cm,angle=0}} \\caption[]{\\label{fg:jetexp} Projected location and size of the emitting feature as derived from the observed spectral evolution. The locations of the emitting feature are calculated assuming a viewing angle of $\\theta_{\\rm j}=3\\fdg2$ obtained from equating the core-to-nozzle travel time with the observed time-lag between the gamma-flare and the ejection of B3. The respective spectral epochs 1--7 are given next to the data points. Open square denotes the location of the VLBI-core of \\object{0836+710}. The increase of the size of the emitting region can be modeled by a conical jet with an opening angle $\\phi_{\\rm j} = 2\\fdg1\\pm0\\fdg1$. } \\end{figure} The derived projected location and size of the emitting region are shown in Figure \\ref{fg:jetexp} for all spectral epochs. The open square refers, in Figure \\ref{fg:jetexp}, to the back-extrapolated ejection epoch of B3 (associated with the passage of B3 through the VLBI core). At earlier epochs, the emitting region was optically thick at frequencies $\\nu\\le 22$\\,GHz, which is reflected in its negative location with respect to the VLBI core (implying that the emitting region was traveling between the jet base and the VLBI core). The region was expanding almost linearly; the resulting opening angle $\\phi_{\\rm j} = 2\\fdg1 \\pm 0\\fdg1$ is in a good agreement with the observed opening angle measured in the immediate vicinity of the VLBI core (O96). During the expansion, the magnetic field decreases from $\\approx 0.2$\\,G to $\\approx 0.05$\\,G, with the core magnetic field, $B_{t=1992.62} \\approx 0.1$\\,G. With the jet parameters determined, we can now discuss the energy balance in the compact jet of \\object{0836+710}. We adopt the formulation of Blandford and K\\\"onigl (1979) to estimate the total (kinetic$+$magnetic field) power of the jet, $L_{\\rm tot} = 2.0\\cdot10^{48} $\\,erg\\,s$^{-1}$. The corresponding synchrotron power of the jet is $L_{\\rm syn} = 2.5\\cdot 10^{47} $\\,erg\\,s$^{-1}$. The data from von Linde et al. (1993) provide an estimate of the energy deposited in the jet by the optical flare, $E_{\\rm opt} \\approx 2\\cdot 10^{47} $\\,erg. The maximum observed gamma luminosity of \\object{0836+710} is $\\approx 1.8\\cdot 10^{48} $\\,erg\\,s$^{-1}$ (Mukherjee et al. 1997), but it should be made smaller by $10^2$--$10^3$ to account for Doppler boosting of the gamma-radiation. Comparing the optical energy output and gamma luminosity to the jet radio power, we can conclude that flaring activity is not likely to play a significant role in forming and maintaining the jet. On the other hand, we find that the average isotropic luminosity of the optical flare is similar to the isotropic radio luminosity observed at the epoch of the flare ($\\approx 4\\cdot 10^{45}$\\,erg\\,s$^{-1}$), suggesting the relation between the emission in the two bands. In our view, this supports the schemes in which relativistic electrons (responsible for synchrotron emission in the radio and optical bands) are the primary radiating particles in the jet, and external photons (producing the observed gamma-emission) are Compton scattered by the jet electrons (e.g. Blandford \\& Levinson 1995). Then the energy densities of the synchrotron and Compton emission can be related through the Doppler factor, so that $u_{\\rm C}/u_{\\rm syn} = k_e \\delta_{\\rm eq}^{1-\\alpha_{\\rm C}}$, (Dermer, Sturner \\& Schlickeiser 1997). We assume equipartition ($k_e = 1$) between the external radiation field and the magnetic field of the emitting region. Then, for a typical spectral index of Compton emission, $\\alpha_{\\rm C}=-1.0$, we obtain $\\delta_{\\rm eq} = 15.3^{+1.5}_{-1.1}$, which is consistent with the $\\delta_{\\rm j}=16.4$ derived in section \\ref{sc:travel}." + }, + "9802/astro-ph9802302_arXiv.txt": { + "abstract": "We compare the evolution of morphology in X--ray cluster data and in clusters obtained in a simulation of flat CDM, normalized to the observed cluster abundance. We find that the evolution rate in model clusters is significantly higher than in data clusters. The results is stricking, as our cluster data are all just for $z \\simlt 0.2$, but is not contrary to expectations, as is known that the CDM spectrum is too flat, around the cluster scale, and therefore induces the presence of too many substructures. The cluster simulations were run using a TREESPH code and therefore include hydrodynamical effects. The test we performed, which turns out to be very sensitive, starts from the so--called power ratios introduced by Buote and Tsai, and makes use of the 2--dimensional generalization of the Kolmogorov--Smirnov test. The discrepancy between data and model is however directly visible through linear regression techniques and using simpler statistical tests. ", + "introduction": "The principal aim of this work is to study the evolution of morphology in hydrodynamical simulations of galaxy clusters, for flat CDM cosmological models, normalized in order to reproduce the observed cluster abundance. We also compare its trend with observational X--ray data and find a substantial discrepancy on the rate of evolution of substructures, which seem to evolve faster in models than in data. This result is not surprising, as the slope of CDM spectrum, around the cluster scale, is known to be too flat. E.g., the {\\sl extra power} parameter $\\Gamma = 7.13 \\cdot 10^{-3}\\left( \\sigma_8 / \\sigma_{25} \\right)^{10/3}$ ($\\sigma_{8,25}$ are mass variances on $8,25 ~ h^{-1}$Mpc scales; $h=H/100~ {\\rm km}\\, {\\rm s}^{-1}{\\rm Mpc}^{-1}$) is significantly lower in data than in model predictions: for APM galaxies, Peacock and Dodds (1995) found $\\Gamma =0.23 \\pm 0.04$; for the Abell/ACO sample, Borgani \\etal (1997) found $\\Gamma$ in the interval 0.18--0.25; on the contrary CDM models predict $\\Gamma \\simeq 0.4\\, $. What is however stricking is the sensitivity that the tests we performed on hydrodynamical simulations, seem to have in appreciating the spectral slope. High $\\Gamma$ values imply that too many \\sub ~intervene in the dynamical growth of fluctuations, on cluster scales. An excess evolution can be therefore expected. But available data cover only a narrow redshift interval ($0 < z < 0.2$) and simulation outputs shall be used on the same interval, reproducing data distribution. The point is then that different rates seem already visible on such redshifts, which concern a fairly recent cosmic evolution. The tests we use start from the a statistical tool introduced by Buote \\& Tsai (1995), who defined the so--called power ratios \\apim, which essentially derive from a multipole expansion of X--ray surface brightness. In sec. 2 more details on \\apim definition will be reported. Much work on \\apim was already done, both in order to derive them from data and to compare them with simulations (Buote \\& Tsai 1995, Buote \\& Tsai 1996, Tsai \\& Buote 1996, Buote \\& Xu 1997). In a recent work, Buote e Xu (1997) gave some plots where the evolution of cluster morphology in data and CDM simulations are compared; similar plots are given also by us. Here, however, we present two essential improvements: (i) Instead of using pure CDM N--body simulations, we use the outputs of a TREESPH code, therefore including hydrodynamical effects. A large number of massive clusters are treated in this way and hydrodynamical effects are found to be important in shaping the distribution of the X--ray emitting baryon component. This will be shown in some plots where DM and baryon distributions are compared; but there is also a back effect of the different baryon distributions on the evolution of the potential well, which cannot be appreciated directly from the plots. (ii) Furthermore, to compare data and simulations, we use an advanced statistical discriminator (Peacock 1982, Fasano \\& Franceschini 1987), which allows us to find quantitative estimates of the probability that the same process gives rise to observational and simulated statistical outputs. The discrepancy we find is however directly visible in suitable evolutionary plots and can be also appreciated using a simpler statistical test. The plan of the paper is as follows. In Sec. 2 power ratios are defined and their evolution is discussed. In Sec. 3 we give some details about the model and simulation procedure followed to define clusters and study their evolution. Power ratios are computed in Sec. 4 where we also perform the comparison between data and simulations. Results are discussed in Sec. 5, where some general conclusions are also drawn. Appendix A includes a brief outline of the PF2 statistical test, and some discussion of the reasons why its extension to three dimensions is premature. ", + "conclusions": "It has been known for several years that COBE normalized CDM yields a cluster abundance exceeding observation by a factor $\\sim 20$. In order to obtain a fair number of clusters, the quadrupole data has to be lowered down to $\\sim 9\\, \\mu$K, at more than 3$\\, \\sigma$'s from the observational value. However, CDM models are still currently in use as reference models; in fact, 0CDM models, mixed models or $\\Lambda$ models, used to obtain fair fits of observational data on various scales, involve more parameters. According to our results, however, even though CDM is normalized to yield a fair number of clusters, we cannot expect their morphologies to be adequately approached by the model. This output is related to the slope of the transfered spectrum, as already outlined in Sec.~1, and as is expressed by the values of the {\\sl extra power} parameter $\\Gamma$. It is however important and fairly unexpected that the effects of such different slope are already visible in the average evolutionary tracks, for a fairly narrow redshift range ($02$) and are evoloving passively or most of them formed at $z<0.5$, as implied by the deficit of S0s in intermediate redshift ($z\\sim0.5$) clusters. The resolution of this controversy may be that the apparent deficit of S0s has been derived from a quantity -- the E to S0 ratio -- which is prone to morphological classification errors. Once all sources of error are taken into account, the E to S0 ratios of clusters at different redshifts are fully compatible, and no additional creation of S0s at $z<0.5$ is required by the data. Furthermore, there is no deficit at all of S0s in the intermediate redshift cluster for which we have morphological types of very high quality, and thus derive an E to S0 ratio with a small error. ", + "introduction": "The S0 class is often considered a class of transition between Es and Ss (see, for example, Hubble 1936, but see van den Bergh 1976 for an opposite opinion). Evidence for the evolution of the S0s was cast looking for an heterogeneity of this class, heterogeneity that suggests a variety of evolutionary paths which end in a S0 galaxy. Color and color dispersion of S0s (Sandage, \\& Visvanathan 1978; Bothun, \\& Gregg 1990; Bower, Lucey, \\& Ellis 1992), segregation in clusters (Oemler 1974; Dressler 1980b), location in the fundamental plane (Saglia, Bender, \\& Dressler 1993), geometrical properties (Michard 1994), the presence of unusual features such as shells, ripples (Schweitzer et al., 1990): all these factors give contradictory indications on the origin of S0s galaxies. Recently, the advent of the {\\it Space Telescope} has permitted a major breakthrough: its superb angular resolution allows the morphological classification of distant galaxies and therefore the comparison of the properties of the morphological types at quite different redshift, i.e. at different look-back times. The evolutionary history of the morphological types should no more be indirectly inferred by the relics it leaves in the galaxy properties, but can be caught in the act. Two recent papers by the MORPHS collaboration (Smail et al. 1997; Dressler et al. 1997) present evidences for the evolution of the S0s class between $z\\sim0.5$ and $z\\sim0$. MORPHS noted that in intermediate redshift ($z\\sim0.5$) clusters there is a deficit of S0s with respect to nearby clusters. This situation is a quite puzzling: intermediate redshift lenticulars are red and with little scatter in color, and therefore are thought to be as old as ellipticals (Ellis et al. 1996; Andreon, Davoust, \\& Heim 1997). Nevertheless their number seems very low, so that many of them must be created in order to reach the population of S0 in nearby clusters (Dressler et al. 1997). Furthermore, at least Coma lenticulars are too old and homogeneous in their properties to be formed at $z<0.5$ (Bower et al. 1992; Andreon 1996). Therefore, S0s in clusters should have $z_{formation}<0.5$ and $z_{formation}>2$ at the same time. The deficit of S0s is not shared by all clusters. In one of the three clusters at $z\\sim0.3$ studied by Couch et al. (1998) the S0s fraction is quite similar to that found in nearby clusters, whereas in the two other it is intermediate between nearby and $z\\sim0.5$ clusters. In the cluster Cl0939+4713 ($z\\sim0.4$) S0s are as frequent as in the nearby Coma cluster (Andreon et al. 1997a), a result which would imply no evolution with time for the S0 population. The situation is striking for this clusters since Dressler et al. (1997) and Andreon et al. (1997a) find opposite results for the same sample of galaxies and from very similar analysis of the same WFPC2 post-refurbished images of the {\\it Hubble Space Telescope}. Therefore, a further analysis is useful in order to understand these discrepancies. ", + "conclusions": "Smail et al. (1997) and Dressler et al. (1997) are aware that the claimed deficit of S0s in intermediate redshift clusters holds only in the absence of redshift dependent errors in their morphological classification. For quantifying systematic errors, they compared the ellipticity distribution of the morphological types at intermediate redshift and of the nearby universe. As nearby sample they used the Revised Shapley-Ames catalogue (Sandage, Freeman, \\& Stokes 1970) and a magnitude complete sample of galaxies in Coma (Andreon et al. 1996), the latter being preferable (Smail et al. 1997) because the ellipticities were measured approximatively at the same isophote of the compared sample. MORPHS finds that the ellipticity distribution of intermediate redshift and nearby Hubble classes are compatible, and conclude that, in spite of the 20 \\% of scatter in the morphological types present in their sample, their classification is equally good and with no bias with redshift. However, this comparison is intended to look for redshift dependent trends in the morphological classification by means of an indirect quantity, the ellipticity distribution of the Hubble types, and does not test directly the quantity of interest, the redshift dependence of the morphological classification, as we do in our comparisons for Coma and Cl0939+4713 morphological types. To summarize, we have shown that the E to S0 ratios of nearby and intermediate redshift clusters are equal within the errors once the error on the morphological type, which is the dominant term, is taken into account. The comparison of equally good morphological types measured by means of independent morphological schemes of the galaxies in CL0939+4713 confirms for this cluster that the claimed deficit of S0s is the result of having measured it by a quantity too prone to morphological errors. Therefore the evolution of the S0 fraction from $z\\sim0.5$ to the present time is not longer required by the MORPHS data (or better by the present analysis of the data), thus solving the puzzle of the old age of the S0s and their absence at $z\\sim0.5$. Furthermore, there is no deficit at all of S0s in the intermediate redshift cluster for which we dispose of high quality morphological types and thus of a E to S0 ratio with small error. However, our conclusion should not be overinterpreted. First of all, we have not shown the absence of an evolution in all the properties of early type galaxies, from $z\\sim0.5$ to the present time. We have just shown that there is no evidence for an evolution of the relative fraction of early type galaxies. Strickly speaking, the constancy of the E to S0 ratio does not exclude morphological changes of individual galaxies between these two classes, and from S to E or S0 classes, but just gives some constraints on the evolution of the two populations. Individual galaxies can change their morphological type while keeping the E to S0 ratio constant, provided that the same number of Es become S0s as viceversa, or, when spiral galaxies are involved, that S become E or S0 with the right frequency. Furthermore, a 20\\% of scatter (error) in the morphological type, claimed both by MORPHS and by ourselves (this work and Andreon, \\& Davoust 1997) is not a negative judgement of the MORPHS work, since it is usual in all morphological studies -- so usual that better agreements are suspicious (Andreon, \\& Davoust 1997)--. The presence of systematic errors in the morphological classifications does not make MORPHS types useless, since they are still useful for studying quantities less affected by systematic errors. A better morphological type for all intermediate galaxies, such as a structural one, would need a 60 time larger effort (Andreon \\& Davoust 1997) and thus is out of our present capabilities. Recently, Van Dokkum et al. (1998) found that at large clustercentric radii ($R>0.7 h^{-1}_{50}$ Mpc) of the intermediate redshift cluster CL1358+62, S0s are heterogeneous in color and therefore experienced star formation until very recently, reaching conclusions opposite to those from previous works focussed on the central regions of other intermediate redshift clusters (Ellis et al. 1996; Andreon, Davoust, \\& Heim 1997). They suggest that S0s evolve primally in the transition region between the cluster and the field, giving support to MORPHS finding that S0s are still forming at intermediate redshift. However, van Dokkum et al. (1998) choose to classify as S only starforming galaxies, leaving spirals with faint smooth spiral structure in the S0 class (see their Section 2.3.1). This is confirmed by the inspection of their black and white prints presented in their Figure 3: at least 25 \\% of all galaxies that van Dokkum et al. classify as S0s are S according to MORPHS or our morphological scheme. All these galaxies present smooth spiral arms (or irregular isophotes) and look as the spirals in CL0939+4713 shown in Figure 3 or to Coma spirals. Therefore, Van Dokkum et al. (1998) adopt a classification scheme different from the one adopted by other researchers. The shape of the azimuthal averaged surface brightness profile measured by van Dokkum et al. (1998) does not discriminate S0s from early--type spirals, since the difference between the two types is given by the smooth spiral arms whose contribution to the radial surface brightness profile is negligible. Their conclusion about the evolutive nature of S0s is therefore relative to their S0 class, and not to the Hubble S0 class. A reclassification of van Dokkum et al. S0s in the Hubble scheme would help to understand which part of the heterogeneity of their S0s class is due to the pollution by early--type spirals and which part is intrinsic and suggestive of recent formation. At small clustercentric radii, where the misclassification is likely low due to the morphological segregation, van Dokkum et al. S0s are an homogeneous old population of galaxies, as in all the other studied clusters both in the nearby universe and at intermediate redshift. As a final remark, we stress that the morphological type is a quantity which needs to be calibrated and with an associate error, as all other physical quantities. We re-iterate the need of using a classification method which keep minimal morphological type errors and systematic differences from the standard scale: the Hubble sequence." + }, + "9802/physics9802019_arXiv.txt": { + "abstract": "We discuss the generation and statistics of the density fluctuations in highly compressible polytropic turbulence, based on a simple model and one-dimensional numerical simulations. Observing that density structures tend to form in a hierarchical manner, we assume that density fluctuations follow a random multiplicative process. When the polytropic exponent $\\gamma$ is equal to unity, the local Mach number is independent of the density, and our assumption leads us to expect that the probability density function (PDF) of the density field is a lognormal. This isothermal case is found to be singular, with a dispersion $\\sigma_s^2$ which scales like the square turbulent Mach number $\\tilde M^2$, where $s\\equiv \\ln \\rho$ and $\\rho$ is the fluid density. This leads to much higher fluctuations than those due to shock jump relations. Extrapolating the model to the case $\\gamma \\not =1$, we find that, as the Mach number becomes large, the density PDF is expected to asymptotically approach a power-law regime, at high densities when $\\gamma<1$, and at low densities when $\\gamma>1$. This effect can be traced back to the fact that the pressure term in the momentum equation varies exponentially with $s$, thus opposing the growth of fluctuations on one side of the PDF, while being negligible on the other side. This also causes the dispersion $\\sigma_s^2$ to grow more slowly than $\\tilde M^2$ when $\\gamma\\not=1$. In view of these results, we suggest that Burgers flow is a singular case not approached by the high-$\\tilde M$ limit, with a PDF that develops power laws on both sides. ", + "introduction": "The formation of density structures by the velocity field of highly compressible turbulence is of great interest in astrophysics. The determination of their typical amplitude, size and volume filling factor poses significant difficulties since it requires a knowledge of the full statistics. As a first step, we shall concentrate in this paper on one-point statistics and more specifically on the probability density function (PDF) of the density fluctuations in one-dimensional (1D) turbulent flows. It is well known that the density jump in a shock depends directly on the cooling ability of the fluid. Thus, for an adiabatic flow the maximum density jump is 4, for an isothermal flow it is $\\sim M_a^2$ \\cite{landau}, and for nearly isobaric flows it is $\\sim e^{M_a^2}$ \\cite{VPP96}, where $M_a$ is the Mach number ahead of the shock. The net cooling ability of a flow can be conveniently parameterized by the polytropic exponent $\\gamma$, so that the thermal pressure $P$ is given by $P=K\\rho^\\gamma$, where $\\rho$ is the fluid density \\cite{footnote}. Isothermal flows have $\\gamma=1$, and isobaric flows have $\\gamma=0$. Note that $\\gamma<0$ corresponds to the isobaric mode of the thermal instability (see, e.g., \\cite{B95}). Thus, in general, the amplitude of the turbulent density fluctuations will be a function of $\\gamma$. Previous work with isothermal flows had suggested that the PDF is log-normal \\cite{V94,PNJ97}, while for Burgers flows a power-law PDF has been reported \\cite{GK93}. More recently, evidence that flows with effective polytropic indices $0<\\gamma<1$ also develop power-law tails at high densities has been presented \\cite{SVCP97}. In order to resolve this discrepancy, we present a series of 1D numerical simulations of polytropic gas turbulence with random forcing, in which the polytropic exponent $\\gamma$ parameterizes the compressibility of the flow. We have chosen to use 1D simulations in order to perform a large number of experiments at a sufficiently high resolution, integrated over very long time intervals, allowing us to collect large statistical samples. The simulations have three governing parameters: the polytropic index $\\gamma$, the Mach number $M$, and the Reynolds number $R$. We keep the Reynolds number fixed, and explore the effects of varying $\\gamma$ and $M$ on the resulting density PDF. We find that varying these two parameters is not equivalent. Variation of $\\gamma$ induces a clear qualitative variation of the density PDF, which, at large Mach number, displays a power-law tail at high densities for $0<\\gamma<1$, becomes log-normal at $\\gamma=1$, and develops a power-law tail at low densities for $\\gamma>1$. This suggests a symmetry about the case $\\gamma=1$, which we also explore. Variation of the Mach number, on the other hand, only appears to induce a quantitative change, in such a way that increasing $M$ augments the width of the PDF. The plan of the paper is as follows. In sec.\\ \\ref{numerics} we describe the equations solved and the numerical method. In sec.\\ \\ref{statistics} we describe the statistics of the various fields, in terms of their PDFs, together with a tentative model and a discussion of the Burgers case. Section \\ref{conclusions} is devoted to a discussion on the choice of the forcing, together with a summary of our results. ", + "conclusions": "\\label{conclusions} \\subsection{Effects of the forcing} The study presented in this paper has been performed for a single choice of the forcing and of the Reynolds number. While the variation with the latter parameter can be trivially extrapolated, we cannot a priori be sure that our results are independent of the type of forcing. We have performed decay runs and observed that the behavior of $\\sigma_s$ vs. $\\tilde M$ is still the same as in the forced case. The PDFs however cannot be computed on a single snapshot due to the poor statistics and cannot be integrated in time since the Mach number changes by roughly one or two orders of magnitude during the run. We have also performed a run at $\\gamma=1$ with a forcing of the form $f/\\rho$ in eq. (\\ref{eq:baseu}). In that case the density PDF is not a lognormal anymore but presents a power law tail for low densities (not shown). This can be attributed to the fact that the flow is stirred more vigorously at low densities so that the effective Mach number indeed increases as $\\rho$ decreases. We nevertheless think that our results can be extrapolated to an unforced situation, at a given time, and possibly also to the multi-dimensional case. Note that the Mach numbers we have explored in this paper would correspond to even higher Mach numbers in the multi-dimensional case since in that case only a fraction of the total kinetic energy populates the compressible modes. \\subsection{Summary} We have presented an investigation of the density PDFs of a randomly accelerated polytropic gas for different values of the polytropic index and of the Mach number. We have suggested a simple model in which the density field is everywhere constructed by a random succession of jumps \\cite{V94}. When the flow is isothermal ($\\gamma=1$), the jumps are independent of the initial density, and have always the same probability distribution. Expressed with the variable $s\\equiv \\ln \\rho$ the jumps are additive, and by the Central Limit Theorem are expected to have a Gaussian PDF, or a lognormal in $\\rho$. An analysis of the expected $s$ increments in the weak and strong shock cases, as well as those due to expansion waves, suggested that the variance $\\sigma_s^2$ should scale as the mean square turbulent Mach number $\\tilde M^2 $. Moreover, because of mass conservation, the peak of the distribution $s_o$ is related to the variance by $s_o=-\\frac{1}{2}\\sigma_s^2$. These predictions were verified in 1D simulations of compressible turbulence. Previous claims that it is the {\\it density} variance $\\sigma_{\\rho}^2$ that should scale as $\\tilde M^2$ \\cite{PNJ97} might have been misled by lower effective Mach numbers than those achieved in the present simulations, in which all of the kinetic energy is in compressible modes thanks to the one-dimensionality. When $\\gamma \\not=1$, the density jumps are not independent of the local density anymore, and the shape of the PDF should change. Observing that a renormalization of the Mach parameter (eq.\\ (\\ref{eq:baseu})) $M\\rightarrow M(s;\\gamma)= Me^{\\frac{1-\\gamma}{2}s}$ restores the form of the equations for the case $\\gamma=1$, we proposed the ansatz that the PDF may still be described by the same functional form as in the case $\\gamma=1$, but substituting $M$ by $M(s;\\gamma)$. This prediction is confirmed by the numerical simulations, giving PDFs which are qualitatively in very good agreement with the model PDF, eq.\\ (\\ref{eq:PDFgne1}). The result is that the PDF asymptotically approaches a power law on the side where $(\\gamma-1)s<0$, while it decays faster than lognormally on the other side. Upon the replacements $\\gamma \\rightarrow (2-\\gamma)$ and $\\rho \\rightarrow 1/\\rho$ we find, using the condition of mass conservation, that the slope $\\alpha$ of the power law for a given value of $\\gamma$ is related to its value at $2-\\gamma$ by eq.\\ (\\ref{eq:symalpha}) in the large Mach number limit. These results are also confirmed by the numerical simulations, which exhibit a power law at $s>0$ when $\\gamma<1$ and at $s<0$ when $\\gamma>1$, with slopes which are roughly related by eq.\\ (\\ref{eq:symalpha}), with better accuracy at large Mach numbers. Finally, on the basis of these results, we suggested that the Burgers case should develop a power law PDF at both large and small densities, since in this case there is no pressure on either side. This result was again confirmed by a simulation of a Burgers flow. We shall conclude this paper by pointing out that the non-uniqueness of the infinite Mach number limit might have important consequences for astrophysical applications, such as in cosmology. The so-called Zeldovich \\cite{Z70} approximation is indeed based on the Burgers' equation which, at the light of the present work, appears as a questionable model of highly compressible flows. This point will be addressed in future work." + }, + "9802/astro-ph9802354_arXiv.txt": { + "abstract": "We show that a 2D projection is representative of its corresponding 3D distribution at a confidence level of 90 $\\%$ if it follows a King profile and if we consider the whole spatial distribution. The level is significantly lower and not decisive in the vicinity of the 2D cluster center. On another hand, if we verify the reciprocal statement of the Mattig's distribution (1958) -i.e. a flux limited sample is represented by a 0.6 slope of its count law-, we point out that, due to the usual unaccuracy of the slope determination, a slope of 0.6 is not a sufficiently strict criterion for completeness and uniformity of a sample as often used in the literature. ", + "introduction": "A large part of modern observational cosmology focuses on statistical studies of the galaxy population in the Universe. Two problems are commonly faced that concern the properties of galaxy samples: First, the spatial distribution of galaxies belonging to clusters is basically unknown. The image of a cluster is a 2-dimensional projection of the true 3-dimensional distribution, and one is led to question the quality of this 2D projected distribution as estimator of the true 3D one. Let us recall that this point is important for understanding the true physics of the cluster, and that properties as the core radius (Lubin \\& Postman 1996) and morphological segregation (e.g. Whitmore \\& Gilmore 1991, Stein 1996) are derived from the 2D distributions only. Second, one usually uses the slope of the count law of flux-limited extragalactic samples to estimate their completeness (e.g. Paturel et al. 1994) and uniformity. This test is based on Mattig's demonstration (1958) that, in an Euclidian Universe, the number $N$ of objects brighter than an $m$ magnitude is: $N(m) \\propto 10^{0.6m}$. However, the validity of the reciprocal statement, that a count law with slope 0.6 in ($m, log N$) coordinates implies uniformity and completeness, remains to be checked carefully. The present paper aims at clarifying these two questions by means of Monte-Carlo simulations based on input parameters consistent with observed ones. For this purpose, a serie of test sample, containing either simulated and real data have been considered. For this purpose, series of test samples containing either simulated or real data, have been considered. The samples are presented in Section II. Section III describes the relationship between the projected 2D and real 3D distributions of galaxies in clusters. Section IV discusses the slope of counts laws and completeness. Section V contains our conclusion and summary. ", + "conclusions": "We have shown that a 2D projection is representative at a level of 90$\\%$ of the corresponding 3D sample if it obeys a King profile and if we consider the whole cluster. The confidence level is significantly lower in the vicinity of the cluster 2D center (typically inside 4 core radii) and is not decisive there. We have also shown that it is somewhat unsafe to use the slope of the count law of a sample to test its completeness level and its uniformity. Indeed, for the uniformity of the sample, we have almost no observable difference between a realistic artificial sample (with 2/3 of its population in clusters), and a completely uniform artificial sample. For the completeness level, the difference is about 15\\% between a complete sample and an incomplete one. This difference represents an error lower than the usual accuracy of the slope determination. We conclude that the slope of the count law, in itself, is not a decisive factor to assess uncompleteness and homogeneity of a sample." + }, + "9802/astro-ph9802162_arXiv.txt": { + "abstract": "Enigmatic transitions between spin-up and spin-down have been observed in several \\Xray\\ pulsars accreting matter via an accretion disk. In these transitions, the torque changes sign but remains at nearly the same magnitude. It has been noted previously that alternating prograde and retrograde disk flows would explain many features of the torque reversals, although it has been unclear how a stable retrograde disk could be formed. We suggest that the reversals may be related to the disk at times being warped to such an extent that the inner region becomes tilted by more than 90 degrees. This region would thus become retrograde, leading to a negative torque. Accretion disk models can show such behavior, if account is taken of a warping instability due to irradiation. The resulting `flip-overs' of the inner parts of the disk can reproduce most characteristics of the observations, although it remains unclear what sets the timescale on which the phenomenon occurs. If this model were correct, it would have a number of ramifications, for instance that in the spin-down state the \\Xray\\ source would mostly be observed through the accretion disk. ", + "introduction": "\\label{sec:intro} Long-term, continuous monitoring by the BATSE all-sky monitor on the {\\it{}Compton Gamma Ray Observatory} has led to a qualitative change in our picture of the spin-frequency behavior of accreting \\Xray\\ pulsars (Bildsten et al.\\ 1997, hereafter \\cite{bild&a:97}). Of the four well-measured persistent sources thought to accrete by way of an accretion disk, all display sudden transitions between episodes of steady spin-up and spin-down. The spin change rate (i.e., the absolute value of the pulse frequency derivative) is nearly equal for spin-up and spin-down, and it is comparable to the torque expected if all the angular momentum of the accreting gas is deposited at the magnetospheric boundary and transferred to the neutron star. In at least some of these systems, however, Roche-lobe overflow is thought to occur, and hence the matter inserted into the accretion disk has only one sense of angular momentum; symmetric torque reversals would thus not be expected. Previous accretion torque models, in which the net torque could be positive or negative depending on the mass accretion rate $\\dot M$ (e.g., \\cite{ghosl:79}), cannot easily reproduce the observations (see \\cite{nels&a:97}). We only list the main arguments here: (i) one needs step-wise changes in $\\dot{M}$ to produce distinct spin-up and spin-down states -- this seems unlikely; (ii) one would expect changes in $\\dot{M}$ to be reflected in the \\Xray\\ luminosity $L\\sub{X}$ -- there are variations in $L\\sub{X}$, but these appear uncorrelated with spin-up/down state; and (iii) one expects at all times a positive correlation between torque and luminosity -- for \\GX\\ an anti-correlation is observed during its spin-down state (\\cite{chak&a:97b}). Suggestions about a possible cause for the torque reversals have been made by Yi, Wheeler \\& Vishniac (\\ctyr{yiwv:97}) and Nelson et al.\\ (\\ctyr{nels&a:97}). Yi et al.\\ suggested that the reversals were due to small changes in $\\dot{M}$ around a critical value at which the system changes from a primarily Keplerian flow to a substantially sub-Keplerian, radially advective flow. This addresses points (i) and (ii), but not point (iii). Nelson et al.\\ (\\ctyr{nels&a:97}) explored the possibility of having systems in which nothing changes except the sense of rotation of the disks. They found that this would explain the observations very well. They quoted a suggestion by Makishima et al.\\ (\\ctyr{maki&a:88}), that in \\GX\\ one might have accretion from a wind and thus form a retrograde disk more easily. As Nelson et al.\\ noted, however, it is very hard to imagine how a stable retrograde disk could form, especially in the ultracompact binary \\fouru, for which all indications are that mass transfer is by Roche lobe overflow from a very low-mass companion. Here, we suggest a modification of the Nelson et al.\\ picture, viz., that it is only the inner part of the disk that is changing its sense of rotation, as a consequence of very strong warping of the disk. Accretion disks are unstable to warping if lit strongly by a central radiation source (\\cite{prin:96}). In numerical simulations which include such irradiation, the disk can sometimes become more and more warped, until the inner part has become inclined by well over 90 degrees (\\cite{prin:97}; see Fig. 6 of that paper for illustrations of a disk with its inner parts flipped over). In Section~\\ref{sec:obs}, we give an updated summary of the observations. In Section~\\ref{sec:warp}, we briefly discuss the warping process and describe simulations done specifically for \\Xray\\ binaries. We proceed to make a qualitative comparison with the observations, and to discuss ramifications. We summarize our conclusions in Section~\\ref{sec:conc}. ", + "conclusions": "\\label{sec:conc} It has been previously suggested that retrograde disk flows would be an elegant explanation for torque reversals in disk-fed X-ray pulsars. It has been an open question, however, how such a flow would be created when mass is transferred by Roche-lobe overflow and thus enters the accretion disk in a prograde orbit. We have pointed out that strong disk warping may produce a retrograde flow close to an accreting neutron star from an initially prograde flow in the outer disk. While this simple picture allows us to understand some of the characteristics of the observations, it does not match all of them. This could be because the models currently neglects many potentially important effects. Our main point here, however, is that quasi-stable retrograde disk flows are physically plausible, and that they may thus be (part of) what causes the spin-down states in persistent disk-fed X-ray pulsars. Despite the uncertainties in the models, it remains true generally that if the spin-down state is due to a flipped-over inner disk, it should be associated with enhanced absorption. One would expect relatively strong absorption edges, especially of iron, which could be looked for with high-resolution spectroscopy. The absorption should vary with time as the inner disk precesses (although the precession is not necessarily very coherent). With more disk surface exposed to the central X-ray source, we also would expect increased scattering and stronger fluorescence lines." + }, + "9802/astro-ph9802212_arXiv.txt": { + "abstract": "Are dwarf spheroidal galaxies dark matter dominated? We present N-body simulations of the interaction between the Milky Way and its closest companion, \\sgr\\ galaxy, constrained by new kinematic, distance and surface density observations detailed in a companion paper. It is shown that there is no possible self-consistent solution to the present existence of \\sgg\\ if its distribution of luminous matter traces the underlying distribution of mass. The luminous component of the dwarf galaxy must therefore be shielded within a small dark matter halo. Though at present we are unable to construct a fully self-consistent model that includes both the stellar and dark matter components, it is shown numerically that it is possible that a pure dark matter model, approximating the dark matter halo deduced for \\sgg\\ from analytical arguments, may indeed survive the Galactic tides. The orbit of \\sgg\\ around the Milky Way is considered, taking into account the perturbative effects of the Magellanic Clouds. It is shown that at the present time, the orbital period must be short, $\\sim 0.7\\Gyr$; the initial orbital period for a $10^9\\msun$ model will have been $\\sim 1\\Gyr$. It is found that a close encounter with the Magellanic Clouds may have occured, though the chances of such an interaction affecting the orbit of \\sgg\\ is negligible. ", + "introduction": "The Sagittarius dwarf galaxy (\\markcite{me94}Ibata, Gilmore \\& Irwin 1994, \\markcite{me95c}1995), the closest satellite galaxy of the Milky Way, provides an ideal laboratory in which the complex interactions that take place during the merging of galaxies may be probed. Motivated by these considerations, much information has now been gained on its kinematics, metallicity and stellar populations; the observational constraints obtained hitherto are reviewed in \\markcite{me97}Ibata \\etal\\ (1997; hereafter referred to as IWGIS). A particularly interesting assertion that results from an analysis of these data is that the sheer existence of \\sgg\\ at the present time is very surprising. Accurate kinematic and distance data, which now sample most of the extent of \\sgr, imply (subject to an assumption scrutinized in section~3 below) that this dwarf galaxy has a short orbital period around the Milky Way, less than $\\sim 1$~Gyr. Previously published numerical experiments of the disruption of this dwarf galaxy (\\markcite{vel95}Velasquez \\& White, 1995; \\markcite{joh95}Johnston \\etal, 1995) showed that it is unlikely to survive more than a few perigee passages. Taking the results of these simulations to their logical conclusion, IWGIS argued that the observed stellar population cannot trace the mass of that dwarf galaxy, as \\sgr\\ would have been destroyed by the Galactic tides long ago. A self-consistent solution to the present existence of the dwarf can then only be found if the requirement that light traces mass is relaxed. Using the simple Jacobi-Roche tidal disruption criterion, IWGIS proposed a solution in which the stellar component of the dwarf galaxy is enveloped in a halo of dark matter, which has a mass profile such that dark matter density at the photometric edge of the dwarf is sufficiently high to impede tidal disruption. To be consistent with the observed low velocity dispersion of the stellar component embedded therein, the core radius of the dark halo would have to extend out to the photometric edge of the system. The dwarf spheroidal companions of the Milky Way have long been suspected to contain large quantities of dense dark matter \\markcite{fab83}(\\eg Faber \\& Lin 1983; \\markcite{irw95}Irwin \\& \\des 1995), so the above conclusion for the specific case of \\sgg\\ is perhaps not surprising; however, the density profile of the dark matter deduced by IWGIS has important implications for the nature of the dark halos and their constituents. Most of the dwarf spheroidals contain stars with a broad range of ages and metallicities, which is unexpected in the simplest explanation for their low mean metallicities --- that chemical evolution was truncated by supernovae-driven winds (\\eg \\markcite{san65}Sandage 1965; \\markcite{dek86}Dekel \\& Silk 1986); this problem may be alleviated with the dark matter halo model proposed by IWGIS, due an enhancement of the escape velocity from the dwarf galaxy. These considerations about the dark matter content of \\sgg\\ have substantial implications for the currently-popular hierarchical clustering picture of structure formation, such as Cold-Dark-Matter dominated cosmologies. A very significant accretion and merging of smaller systems occurs during the evolution of a normal galaxy like the Milky Way; is this still an on-going process? In this paper we aim to examine IWGIS' claims, redoing their approximate analytical calculations with numerical disruption experiments. These simulations will be constrained with all available relevant data. In particular, we will first investigate whether the assumption adopted by IWGIS in determining the orbit of \\sgg\\ holds true, and so establish the best-fit orbit; and secondly, we will attempt to find a self-consistent solution to the present existence of the dwarf galaxy. This provides the necessary detailed analysis to determine the validity of IWGIS' assertion. ", + "conclusions": "Comparison of the observed velocity profile to the simulations presented above indicates that \\sgg\\ has a short period orbit, with radial period $T \\sim 0.7 \\Gyr$. It is found that any reasonable model of the internal structure of the dwarf galaxy, where light traces mass and where $M/L \\simlt 10$, either does not survive the tidal interaction with the Milky Way, or has a minor axis half-mass radius that is inconsistent with observations. Thus, it is not possible to understand the present existence of \\sgr\\ if most of its mass is in the form of stars. However, this problem may be solved if Sagittarius, and by implication, other dwarf spheroidal galaxies, have a radially increasing mass to light ratio, as suggested by IWGIS. This analysis supports the mass to light ratio determination of IWGIS, who found $M/L_{\\rm global} \\sim 100$. This conclusion should still hold if there is internal rotation in the dwarf. Nor should the conclusion be affected significantly by the choice of Galactic potential model. The perturbative effect of the Magellanic Clouds was considered; given current estimates of their kinematics, there is a negligible chance that they could have altered sufficiently the orbit of \\sgg\\ to account for the short period deduced above. Further numerical work is required to construct a fully self-consistent model in which both the stellar and dark matter components are present and can reproduce the observations. \\bigskip We are very grateful to Derek Richardson for kindly letting us use of his {\\tt box\\_tree} code and related plotting routines, and especially for many illuminating email conversations. We also cordially thank G. Fahlman, K. Menon, H. Richer, D. Scott and G. Walker for generous loan of their computer resources. RAI expresses gratitude to the Killam Foundation (Canada) and to the Fullam Award for support. \\vfill\\eject" + }, + "9802/astro-ph9802024_arXiv.txt": { + "abstract": "Visible afterglow counterparts have now been detected for two GRBs (970228 and 970508) but are absent, with $L_{opt}/L{\\gamma}$ ratios at least two orders of magnitude lower, for other GRBs, \\eg 970828. The causes of this variation are unknown. Any correspondence which could be discovered between the \\gray properties of a GRB and its $L_{opt}/L{\\gamma}$ would be useful, both in determining the GRB mechanisms, and in allocating resources for counterpart searches and studies. This paper presents the \\gray spectra of GRB 970228 as measured by the Transient Gamma-Ray Spectrometer and comments on characteristics of this GRB compared to others that do and do not have observable counterparts. ", + "introduction": " ", + "conclusions": "" + }, + "9802/hep-ph9802382_arXiv.txt": { + "abstract": "{We show that ``top-down'' mechanisms of Ultra-High Energy Cosmic Rays which involve heavy relic particle-like objects predict Galactic anisotropy of highest energy cosmic rays at the level of minimum $\\sim 20\\%$. This anisotropy is large enough to be either observed or ruled out in the next generation of experiments. } ", + "introduction": " ", + "conclusions": "" + }, + "9802/astro-ph9802030_arXiv.txt": { + "abstract": "We present 7 mm $\\lambda$ VLBA observations of the compact nonthermal radio source in the Galactic Center, Sgr A*. These observations confirm the hypothesis that the image of Sgr A* is a resolved elliptical Gaussian caused by the scattering of an intervening thermal plasma. The measured major axis of Sgr A* is $0.76 \\pm 0.04$ mas, consistent with the predicted scattering size of $0.67 \\pm 0.03$. We find an axial ratio of $0.73 \\pm 0.10$ and a position angle of $77\\degd 0 \\pm 7\\degd 4$. These results are fully consistent with VLBI observations at longer wavelengths and at 3 mm $\\lambda$. We find no evidence for any additional compact structure to a limit of 35 mJy. The underlying radio source must be smaller than 4.1 AU for a galactocentric distance of 8.5 kpc. This result is consistent with the conclusion that the radio emission from Sgr A* results from synchrotron or cyclo-synchrotron radiation of gas in the vicinity of a black hole with a mass near $10^6 M_{\\sun}$. ", + "introduction": "The radiation from the compact radio source in the Galactic Center, Sgr A*, has been clearly determined to be nonthermal in nature. VLBI observations have found the source to have a brightness temperature in excess of $1.4 \\times 10^{10}$ K and a size less than 1 AU (Rogers \\etal \\markcite{roger94} 1994). The compact source has been interpreted as a black hole undergoing accretion at a very low rate ($\\sim 10^{21}$ to $10^{22} {\\rm\\ g\\ s^{-1}}$), either through advection dominated accretion (Narayan, Yi \\& Mahadevan \\markcite{naray94} 1994) or through spherical accretion which may include a disk at small radii (Melia \\markcite{melia94} 1994). In both models the accreting material is mass lost from nearby massive stars through winds. The radio emission may originate in an infalling spheroid, a disk or a low-power Blandford-K\\\"{o}nigl jet (Falcke, Mannheim \\& Biermann \\markcite{falck93} 1993). These theories are constrained by the broadband spectrum and the compact size of the radio source. Recent dynamical evidence supports the black hole hypothesis. The proper motions of early type stars within 0.01 pc of Sgr A* are greater than 1000 km s$^{-1}$ (Genzel \\etal \\markcite{genze97} 1997). The velocity dispersion as a function of radius suggests a contained mass of $2.6 \\pm 0.35 \\times 10^6 M_{\\sun}$. Further, if the kinetic energies of these stars are in equipartition with that of Sgr A*, then we find that the proper motion of Sgr A* (Backer \\markcite{backe96} 1996) implies a minimum mass of $1.4 \\times 10^4 M_{\\sun}$. If this mass is confined within 1 AU as the VLBI results imply, then a mass density of $3 \\times 10^{19} {\\rm M_{\\sun} pc^{-3} }$ is required. Considerations of the short time scales of stability of a cluster of dark objects with this density lead to the strong conclusion that Sgr A* indeed harbors a massive black hole (Maoz \\markcite{maoz97} 1997). Due to diffractive scattering by electrons in the turbulent interstellar medium near the Galactic Center, the angular size of Sgr A* is broadened (e.g., Yusef-Zadeh \\etal \\markcite{yusef94} 1994 and references therein). The source is elliptical at a position angle of approximately $80\\arcdeg$ with an axial ratio near 0.5. The major axis has a $\\lambda^{2.01 \\pm 0.02}$ dependence that extends from 20 to 0.3 cm. The scattering hypothesis is supported by the discovery of a similar angular broadening and asymmetry in the images of OH masers within 15 arcmin of Sgr A* (van Langevelde \\& Diamond \\markcite{vanla91} 1991). Since the effect of scattering decreases more rapidly with wavelength than angular resolution, at a short enough wavelength the compact source should appear unobscured. Currently VLBI observations at millimeter wavelengths stand in conflict. At 3 mm $\\lambda$ Rogers \\etal \\markcite{roger94} (1994) found an upper limit to the apparent size of 0.2 mas, in agreement with an extrapolation of the scattering law. Observing at 7 mm $\\lambda$ in August 1992, Backer \\etal \\markcite{backe93} (1993) found with 5 stations of the National Radio Astronomy Observatory\\footnote{ The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} Very Long Baseline Array (VLBA) an apparent size of 0.7 mas, also in agreement with the scattering law. However, Krichbaum \\etal \\markcite{krich93} (1993), observing with an array of 4 VLBA stations at 7 mm $\\lambda$ in May 1992, found a size $1.7 \\times 0.7$ mas at a position angle of $-20 \\arcdeg$. They also consider a two component model in which the brighter component has a size of $ 0.7 \\pm 0.1$ mas and the fainter component is at a position angle of $-25\\arcdeg$. In this paper we present 7 mm $\\lambda$ observations of Sgr A* using the full VLBA. The additional baselines in these observations significantly improve the density and extent of $uv$ coverage over the previous experiments. We present a map as well as a study of the visibility data. ", + "conclusions": "Imaging with the full VLBA at 7 mm $\\lambda$ shows that the size and shape of Sgr A* is fully consistent with the predictions of the scattering law derived at lower frequencies. We find a major axis FWHM of $0.76 \\pm 0.04$ mas. The scattering model predicts $0.67 \\pm 0.03$ mas. We also find an axial ratio and a position angle consistent with that previously identified at longer wavelengths and with the Backer \\etal \\markcite{backe93} (1993) results at 7 mm. We find no evidence for structure that is not elliptically symmetric to a limit of 35 mJy. We cannot rule out the existence of a second component in the past: a synchrotron component at the epoch of the Krichbaum \\etal (1993) observations is likely to have decayed significantly by the epoch of our observations. If the intrinsic size adds in quadrature with the scattering size to form the apparent size, the major and minor axes of the intrinsic source must be less than 0.48 mas. For a galactocentric distance of 8.5 kpc, this corresponds to 4.1 AU. We infer a lower limit to the brightness temperature of $4.9 \\times 10^9$ K. This limit is less than the limit of $1.4 \\times 10^{10}$ K found by Backer \\etal \\markcite{backe93} (1993) at 7 mm $\\lambda$ and Rogers \\etal \\markcite{roger94} (1994) at 3 mm $\\lambda$ principally due to the lower flux in our epoch of observation. The morphology and size of Sgr A* did not change from the Backer \\etal \\markcite{backe93} (1993) observations despite a factor of two decrease in the flux. The upper limit to source expansion over this epoch is $0.04 {\\rm\\ mas\\ y^{-1}} \\approx 2 {\\rm ~km~s^{-1}}$. The constant size also supports the hypothesis that the apparent source is scattered radiation from the intrinsic source. An unobscured, homogeneous synchrotron source with constant magnetic field, spectral index and peak frequency would show a 40\\% increase in its angular size with a doubling in flux density (Marscher \\markcite{marsc83} 1983). Only a conspiracy of parameter changes could keep the size from changing significantly. The constant ellipticity as a function of wavelength indicates a coherent magnetic field over 0.02 to $3 \\times 10^{-5}$ pc for a scattering screen near the Galactic Center. Detection of a change in the ellipticity with wavelength would be indicative of a turbulent magnetic field on these scales (e.g., Wilkinson, Narayan \\& Spencer \\markcite{wilki94} 1994). At 7 mm $\\lambda$ the refractive time scale is on the order of 2 years. Hence, future observers might expect to see a change in the ellipticity of Sgr A*. These results are consistent with the hypothesis that the radio emission from Sgr A* results from synchrotron or cyclo-synchrotron radiation of gas in the vicinity of a black hole. Both radial and rotating flow models predict the mass of the black hole on the order of $10^6 M_{\\sun}$ and the accretion rate in the range $10^{21}$ to $10^{22} {\\rm\\ g\\ s^{-1}}$ (Melia \\markcite{melia94} 1994, Narayan \\etal \\markcite{naray95} 1995). The radio emission may originate in a spherical cloud or in a low-power inhomogeneous jet. However, the absence of any external feature at any wavelength argues against the jet model. What will shorter wavelength VLBI observations reveal in Sgr A*? Future global VLBI arrays at 1.3 mm $\\lambda$ will have an angular resolution of 20 $\\mu$as (Wright \\& Bower \\markcite{wrigh97} 1997), which is on the scale of a few Schwarzschild radii for a $10^6 M_{\\sun}$ black hole. At this wavelength scattering will not dominate since the expected scattering size is 27 $\\mu$as. However, recent measurements of the centimeter to submillimeter spectrum indicate the presence of two compact sources on two different size scales (Serabyn \\etal \\markcite{serab97} 1997, Falcke \\etal \\markcite{falck97} 1997). If the intrinsic source responsible for the centimeter to millimeter wavelength spectrum has a spectral turnover near 3 mm $\\lambda$, then this component may be forever invisible to VLBI." + }, + "9802/astro-ph9802206_arXiv.txt": { + "abstract": "We present here the results from an extensive scintillation study of twenty pulsars in the dispersion measure (DM) range $ 3-35 \\ {\\rm pc \\ cm^{-3}}$ carried out using the Ooty Radio Telescope, to investigate the distribution of ionized material in the local interstellar medium (LISM). Our analysis reveals several anomalies in the scattering strength, which suggest that the distribution of scattering material in the solar neighborhood is not uniform. Our model suggests the presence of a low density bubble surrounded by a shell of much higher density fluctuations. We are able to put some constraints on geometrical and scattering properties of such a structure, and find it to be morphologically similar to the local bubble known from other studies. ", + "introduction": "Propagation effects on pulsar signals, such as dispersion and scattering, probe the distribution of thermal plasma in the interstellar medium (ISM). Scattering studies of radio waves from pulsars, enable us to probe the electron density fluctuations in the ISM, which are presumably due to turbulence in the ISM (Rickett 1990). Not much is known about electron densities and their fluctuations in the LISM and in order to investigate this in detail, we have made extensive scintillation observations of twenty nearby pulsars. Reliable and accurate estimates of strength of scattering have been obtained and the results are used to study the distribution of electron density fluctuations in the LISM. ", + "conclusions": "" + }, + "9802/astro-ph9802176_arXiv.txt": { + "abstract": "We compare mid-infrared images of the Antennae galaxies (NGC\\,4038/39) from the Infrared Space Observatory, with optical images from the Hubble Space Telescope. The mid-infrared observations show that the most intense starburst in this colliding system of galaxies takes place in an off-nucleus region that is inconspicuous at optical wavelengths. The analyses of the mid-infrared spectra indicate that the most massive stars are being formed in an optically obscured knot of 50 pc radius, which produces about 15\\% of the total luminosity from the Antennae galaxies between 12.5 \\mic\\ and 18 \\mic. The mid-infrared observations reported here demonstrate that the interpretation of star formation properties in colliding/merging systems based on visible wavelengths alone can be profoundly biased due to dust obscuration. The multiwavelength view of this nearby prototype merging system suggests caution in deriving scenarios of early evolution of high redshift galaxies using only observations in the narrow rest-frame ultraviolet wavelength range. ", + "introduction": "One of the most important recent discoveries in extragalactic astronomy has been the identification of a class of ``infrared luminous galaxies\" (L$_{bol}$ $\\geq$ 10$^{11}$ L$_{\\odot}$), which emit more energy in the infrared (5-500 $\\mu$m) than in all other wavelengths combined (see review by Sanders \\& Mirabel, 1996). The trigger of the intense infrared emission appears to be the interaction/merger of molecular gas-rich spirals. Although the spectrum of the integrated radiation from these galaxies suggests that the bulk of the luminosity arises in regions that are heavily obscured by dust, the actual distribution of the mid-infrared emission with high spatial resolution is poorly known. The Infrared Space Observatory (ISO, Kessler \\etal, 1996) offers unprecedented capabilities with respect to the Infrared Astronomical Satellite (IRAS). In the mid-infrared (5.5$\\mu$m - 16.5 $\\mu$m) the Infrared Space Observatory Camera (ISOCAM) provides an improvement in sensitivity of $\\sim$ 1000, and an increase in spatial resolution by a factor of $\\sim$ 60. Furthermore, observations with arcsec resolution in spectrophotometric mode with Circular Variable Filters (CVFs) permit us to infer the nature of the optically invisible stars that heat the dust and ionize the gas. Here we present for the first time an image of NGC\\,4038/39 (Arp\\,244 = VV\\,245 = `The Antennae') in the LW3 (12 -17 \\mic) filter with 1.5$''$ pixel field of view and compare it with the HST optical image. An extensive account of all the ISOCAM observations will be given in a forthcoming paper by Vigroux et al (1998). The Antennae is a prototype nearby merger system of two late-type spiral disk galaxies with nuclei separated by $\\sim$ 6.4 kpc. At a distance of 20 Mpc the total infrared luminosity measured by IRAS is 10$^{11}$ L$_{\\odot}$, which is about five times the luminosity of the system at visual wavelengths. Molecular gas observations with a resolution of 6$''$ (Stanford \\etal\\ 1990) showed that $\\sim$ 60\\% of the CO(1-0) emission from the overall system (Sanders \\& Mirabel 1985) is concentrated in the two nuclei and in an extended off-nuclear complex where the two disks overlap. The spatial distribution of the CO emission is consistent with the $\\lambda$6 cm and $\\lambda$20 cm radio continuum maps by Hummel \\& van der Hulst (1986). ", + "conclusions": "1) The most intense starburst in the prototype merger NGC 4038/39 takes place in an off-nuclear region that is optically obscured. This confirms the interpretation of the high far-infrared luminosity in this system in terms of star formation. 2) The [Ne~III]/[Ne~II] emission line ratios in the mid-infrared indicate that stars as massive as 60 M$_{\\odot}$ are being formed in that optically obscured region. Therefore, to derive the high mass end of the Initial Mass Functions of the starbursts in luminous interacting galaxies, observations at wavelengths longer than the optical are needed. 3) The mid-infrared observations presented here suggest that about 15\\% of the bolometric luminosity from NGC 4038/39 arises in an off-nuclear region that is $\\leq$ 100 pc in size. 4) The most prominent dust lanes in the optical image appear displaced from the peak distribution of the warm dust and gas traced by the mid-infrared emission. This could be due to different spatial distributions of the warm and cool dust, and/or projection effects in the optical appearance of dark lanes. 5) The infrared image also shows strong mid-infrared emission associated with the optically prominent star forming ring in NGC 4038. 6) The effects of absorption by dust are even more dramatic in the UV. Therefore, images of high redshift galaxies in their rest-frame UV could lead to strong biases in the morphological classification, and therefore in the scenarios of the history of galaxy formation." + }, + "9802/astro-ph9802340_arXiv.txt": { + "abstract": "We present a deep BVrIK multicolor catalog of galaxies in the field of the high redshift ($z=4.7$) quasar BR 1202-0725. Reliable colors have been measured for galaxies selected down to $R=25$. Taking advantage of the wide spectral coverage of the galaxies in the field, we compare the observed colors with those predicted by spectral synthesis models including UV absorption by the intergalactic medium and dust reddening. The choice of the optical filters has been optimized to define a robust multicolor selection of galaxies at $3.8\\leq z\\leq 4.5$. Within this interval the surface density of galaxy candidates with $z\\sim 4$ in this field is 1 arcmin$^{-2}$. Photometric redshifts have been derived for the galaxies in the field with the maximum likelihood analysis using the GISSEL library of $\\sim 10^6$ synthetic spectra. The accuracy of the method used has been discussed and tested using galaxies in the Hubble Deep Field (HDF) with known spectroscopic redshifts and accurate photometry. A peak in the redshift distribution is present at $z\\simeq 0.6$ with relatively few galaxies at $z>1.5$. At variance with brighter surveys ($I<22.5$) there is a tail in the distribution towards high redshifts up to $z\\sim 4$. The luminosity function at $z\\sim 0.6$ shows a steepening for $M_B>-19$. This increase is reminiscent of that found in the most recent estimates of the local luminosity function where a similar volume density is reached about 2 mag fainter. The observed cosmological ultraviolet luminosity density is computed in the overall redshift interval $z=0.3-4.5$ reaching a value $\\sim 2\\times 10^{19}$ W Hz$^{-1}$ Mpc$^{-3}$ at $z\\sim 0.8$. Including recent local estimates it appears that the UV luminosity density changes by a factor $\\sim 2.5$ in the overall redshift interval $z=0.1-4$, not including correction for fainter undetected galaxies. Thus we find that the evidence of a marked maximum in the luminosity density at $z\\sim 1-1.5$ for galaxies with $R\\leq 25$ is weak. We have derived in a homogeneous way, using the GISSEL libraries, the physical parameters connected with the fitted spectral energy distributions. Thanks to this new approach, the problem of the star formation history of the universe is dealt with in a self consistent way taking into account the dust and metallicity distributions derived from the spectrophotometric properties of each galaxy in the sample. The bulk of the blue intermediate redshift population with $z=0.4-1$ mostly consists of very young star-forming galaxies with a median starburst age of the order of a few $10^8$ yr and typical mass in luminous stars $\\sim 2\\times 10^8$ M$_{\\odot}$. The presence of this young population is in contrast with pure luminosity evolutionary model (PLE) based on a single high formation redshift. The cosmological mass in formed stars per unit comoving volume at $z\\sim 3$ is already $\\sim 20$\\% of that formed at $z=0.5$ in our magnitude limited sample. Predictions based on the standard hierarchical clustering models are smaller, although not far from that derived from the observations. ", + "introduction": "Deep photometric and spectroscopic surveys of galaxies performed by means of large ground-based telescopes and the HST are providing a first insight into the cosmological distribution of galaxies of different spectral types up to $z\\sim 5$. The bulk of the contribution to surveys of intermediate depth (B$<24.5$, $I<23$, and $K<20$) comes from galaxies at redshifts $z\\sim 0.5$ with a tail in the distribution extending up to $z\\sim 1.5$ (Songaila et al. 1994, Lilly et al. 1995, Ellis et al. 1996, Cowie et al. 1996). These surveys show substantial evolution in the galaxy properties starting already at $z\\geq 0.5$. The presence of massive star formation is provided by the observed blue colors of galaxies and by the detection of strong emission lines. The recognition of the evolutionary path of the average star formation rate per unit comoving volume is one of the main tools to understand the processes that control the formation and evolution of galaxies in the universe. An increase by a factor of ten in the redshift interval $z=0-1$ has been recently suggested (Lilly et al. 1996, Madau et al. 1996). This of course implies that about half of the stars formed at intermediate redshifts, in general agreement with theoretical expectations of CDM cosmologies. Nevertheless, the same cosmological scenarios predict a fraction $<2$\\% of the present mass density in stars at $z>4$ (Cole et al. 1994). It is therefore at these high redshifts that cosmological scenarios for galaxy formation are more vulnerable to observational constraints. Thus, the search and discovery of galaxies at $z>2$ is a major challenge of recent observational cosmology. Until recently, however, these objects have proven to be quite elusive, beeing detected neither in complete spectroscopic surveys limited to $I\\leq 22.5$, (Lilly et al 1995 CFRS, Cowie et al 1996), nor in field surveys based on the search of highly redshifted Ly$\\alpha$ line emission (Thompson et al 1995) with some relevant exceptions (Macchetto et al. 1993, Warren and M$\\o$ller 1996). Quite recently, improved sensitivities and new observational strategies have led to the discovery of star-forming galaxies up to $z\\sim 4.5$. Deep Ly$\\alpha$ searches have been successfully tuned at particular redshifts, such as those of previously known quasars (Hu and McMahon 1996), radio galaxies (Pascarelle et al 1996) or strong QSO absorption systems (Francis et al 1996). Deep multicolor searches are also proving very effective at detecting field galaxies in relatively wide redshift intervals. The technique based on the color selection of the Lyman break feature in flat spectrum star-forming galaxies has been applied by Steidel and co-workers to look for field galaxies at $2.8 < z < 3.4$, (Steidel and Hamilton 1992, Steidel et al 1995), with a remarkable spectroscopic success rate of $\\simeq 70$\\% (Steidel et al 1996). We have also applied a similar multicolor technique to look for the galaxies responsible for strong Ly$\\alpha$ absorption in quasar spectra. By means of deep BVrI images around the $z=4.695$ quasar BR1202-0725 we discovered a galaxy companion whose redshift was estimated to be in the range $z\\geq 4.4-4.7$ (Fontana et al. 1996) on the basis of sharp color criteria obtained comparing observed colors with that expected by star-forming galaxies once UV absorption by the intergalactic medium has been taken into account (Madau 1995). Narrow band imaging and spectroscopic follow up have confirmed the $z=4.7$ redshift of the galaxy (Hu et al. 1996, Petitjean et al. 1996). The majority of galaxies detected in the Hubble Deep Field (HDF; Williams et al. 1996) are too faint to be investigated by spectroscopy. The only way to derive statistical information about their cosmological distribution is based on color-estimated redshifts (Gwinn \\& Hartwik 1996, Lanzetta et al. 1996, Sawicki et al. 1997). The reliability of the method depends mainly on the number of colors used and on the modeling of the galaxy templates. The possibility to calibrate the color-estimated redshifts with fair samples of spectroscopic galaxies allows better accuracy (Koo 1985, Koo \\& Kron 1992, Connolly et al. 1995). We have applied the ``photometric redshift technique'' to derive the redshift distribution of the complete galaxy sample in the field of the quasar BR1202-072 and the results are discussed in this work. To make the redshift estimates more robust we have extended the spectral coverage in the field adding deep K images ($K\\simeq 21.5$ at 3$\\sigma$ sky limit). Photometric redshifts have been derived with a maximum likelihood analysis over the five broad bands BVrIK using a library of $\\sim 10^6$ synthetic spectra. A calibration of the method has been carried out over about 60 galaxies in the HDF with known spectroscopic redshifts. In Sect. 2 we present the data sample, in Sect. 3 the library of galaxy synthetic spectra and our statistical technique for color-estimated redshifts. In Sect. 4 we describe our multicolor selection of galaxies with $z>3.8$ and in Sect. 5,6 we describe and discuss the relevant results concerning the cosmological evolution of the galaxy properties. ", + "conclusions": "Photometric redshifts have been obtained from the $R'\\leq 25$ galaxies selected in the field of the high redshift ($z=4.7$) quasar BR 1202-0725. The wide spectral coverage obtained from deep BVRIK multicolor photometry has allowed an accurate redshift estimate for each galaxy. This has been obtained comparing the observed colors with those predicted by spectral synthesis models including UV absorption by the IGM and dust reddening. We have discussed the accuracy of the method using a control sample of galaxies in the HDF with spectroscopically confirmed redshifts. The main conclusions can be summarized in the following items. \\begin{itemize} \\item A comparison between spectroscopic and photometric redshifts for a sample of galaxies selected in the Hubble Deep Field has shown that spectral synthesis models (Bruzual \\& Charlot 1997) including UV absorption by the IGM and dust reddening can provide photometric redshifts with comparable accuracy ($\\sigma _z \\leq 0.1$ at $z\\leq 1.5$) as obtained by means of the observed spectral templates of local galaxies. The present approach has the advantage to exploit the information on the spectral and physical properties derived from the GISSEL models for each galaxy in the sample. An IMF shape steeper than a Salpeter law provides unbiased redshift estimates and the following statistical results are given for a Miller-Scalo IMF. \\item The redshift distribution of the $R'\\leq 25$ galaxies is peaked at $z=0.6$ with 16\\% of the sample at $z>1.5$. The derived surface density of the $z\\sim 2.8-3.5$ galaxies having $\\langle M_{B_{AB}}\\rangle=-21$ in our field is 1.6 arcmin$^{-2}$ in agreement with that derived in the HDF (1.5 arcmin$^{-2}$) at about the same magnitude. This corresponds to a comoving volume density of $10^{-3}$ Mpc$^{-3}$ similar to the local density of galaxies with the same luminosities. The derived surface density at $3.5-19$ with respect to the extrapolation derived from brighter redshift surveys. A comoving volume density of $2\\times 10^{-2}$ Mpc$^{-3}$ at $M_{B_{AB}}=-17.5$ is obtained. Comparing with the local luminosity function, a luminosity evolution by about 2 magnitudes is suggested for galaxies with $M_{B_{AB}}>-19$. \\item The bulk of the intermediate redshift population mostly consists of very young star-forming galaxies with a median age $\\leq 10^9$ yr and a small stellar mass $M\\sim 5\\times 10^{8} $ M$_{\\odot}$. In particular, the blue fraction with $B'-I'<1.4$ shows a median age of $2\\times 10^8$ yr and stellar mass $M\\sim 2\\times 10^{8}$ M$_{\\odot}$. Any pure luminosity model (PLE) based on a single formation redshift (usually confined to $z>2$) adopted for all the faint field galaxies appears inconsistent with these small age and mass evaluations. \\item The observed 2800 {\\AA} luminosity density and the associated star formation rate in our sample show an increase to $\\phi \\sim 2\\times 10^{19}$ W Hz$^{-1}$ Mpc$^{-3}$ (or $SFR\\sim 5\\times 10^{-2}$ M$_{\\odot}$ yr$^{-1}$ Mpc$^{-3}$) at $z\\sim 0.8$, i.e. only by a factor 2.5 larger than the local value. At $z>1$ the UV luminosity density and the corresponding SFR decrease to values comparable to the local one. Thus evidence of a marked maximum in the luminosity density and SFR at $z\\sim 1$ appears blurred especially if we consider that an significant corrections for fainter undetected galaxies are expected at $z>1$. A comparison between the average cosmological luminosity density and the corresponding star formation rate at $z=0.4-1$ implies an average $E_{B-V}\\simeq 0.1$, adopting the Calzetti (1997) attenuation law and a Miller-Scalo IMF. Finally, the derived cosmological mass in luminous stars per comoving volume at $z\\sim 3-4$ is $\\sim 20-10$\\% of that formed at $z=0.5$, a value larger although not far from that predicted by the standard hierarchical clustering scenario. \\end{itemize}" + }, + "9802/nucl-th9802054_arXiv.txt": { + "abstract": "We present here the first application of realistic shell model (SM) including coupling between many-particle (quasi-)bound states and the continuum of one-particle scattering states to the spectroscopy of $^{8}$B and to the calculation of astrophysical factors in the reaction $^7\\mbox{Be}(p,\\gamma)^8\\mbox{B}$. ", + "introduction": "The theoretical description of weakly bound exotic nuclei close to the drip-line is one of the most exciting challenges today. What makes this subject both particularly interesting and difficult, is the proximity of the particle continuum implying strong modification of the effective nucleon--nucleon interaction and causing unusual spatial properties of the nucleon density distribution (halo structures, large diffusivity). Many of those nuclei are involved in the chain of thermonuclear reactions and, in the absence of data at relevant energies, the models of stars rely to certain extent on calculated astrophysical factors (see Bahcall (1989))~. In weakly bound exotic systems, the number of excited bound states or narrow resonances is small and, moreover, they couple strongly to the particle continuum. Hence, these systems should be described in the quantum open system formalism which does not artificially separate the subspaces of (quasi-) bound (the $Q$-subspace) and scattering (the $P$-subspace) states. For well bound nuclei close to the $\\beta$-stability line, microscopic description of states in the first subspace is given by nuclear SM with model-space dependent effective two-body interactions, whereas the latter subspace is treated in terms of coupled channels equations. In this work, we question the validity of this basic paradigm of nuclear physics, and propose its modification for weakly stable exotic nuclei by taking into account coupling between $Q$~ and $P$~ subspaces in terms of residual nucleon--nucleon interaction. This coupling modifies the scattering solutions as well as the spectroscopic quantities for interior bound states. As said before, we are interested in describing low lying bound and quasi-bound states in exotic nuclei. For that reason, we can restrict description of particle continuum to the subset of one nucleon decay channels. Still, in few rare cases of two-nucleon halo nuclei, this limitation may turn out to be restrictive. In any case, further improvement of our model to more complicated channels like, e.g., $\\alpha$ - channels, can be done as well (Balashov \\etal (1964)). ", + "conclusions": "In this work we have shown results of first calculations using the SMEC which couples the realistic SM solutions for (quasi-) bound states with the scattering solutions of one-particle continuum. The application to $^{7}\\mbox{Be}(p,\\gamma)^{8}$B reaction yields satisfactory description of different components of the radiative capture cross section, including the resonant components. In future, more unstable nuclei should be studied in the SMEC approach to systematically address the problem of effective interactions in the extreme conditions of exotic nuclei. At present, we are applying the SMEC to many other reactions of astrophysical interest such as $^{14}\\mbox{C}(n,\\gamma)^{15}\\mbox{C}$, $^{16}\\mbox{O}(p,\\gamma)^{17}\\mbox{F}$, $^{18}\\mbox{O}(n,\\gamma)^{19}\\mbox{O}$. \\ack Authors would especially like to thank S. Dro\\.zd\\.z and I. Rotter for many clarifying discussions. We are grateful to E. Caurier and A. Lefebvre for encouragement and useful discussions. The work was partly supported by KBN Grant No. 2 P03 B 14010 and the Grant No. 6044 of the French - Polish Cooperation. \\vfill \\newpage" + }, + "9802/astro-ph9802083_arXiv.txt": { + "abstract": "The space-like hypersurface of the Universe at the present cosmological time is a three-dimensional manifold. A non-trivial global topology of this space-like hypersurface would imply that the apparently observable universe (the sphere of particle horizon radius) could contain several images of the single, physical Universe. Recent three-dimensional techniques for constraining and/or detecting this topology are reviewed. Initial applications of these techniques using X-ray bright clusters of galaxies and quasars imply (weak) candidates for a non-trivial topology. ", + "introduction": "If the physical Universe is smaller than the ``observable Universe'', i.e., if the fundamental polyhedron of the Universe is smaller than the sphere of horizon radius in the universal covering space, then some (or many) regions of space will be observable at several (or many) different ``look-back'' times (\\cite{deSitt17}~1917; \\cite{Lemait58}~1958). The word ``Universe'' can be taken here to refer either theoretically to the space-like hypersurface at the present cosmological time, or observationally to the observed past time cone considered in comoving coordinates. Since space-like hypersurfaces are three-dimensional, use of three-dimensional information on astrophysical objects known to exist in the covering space provides a straight-forward way to search for or constrain the global topology of the Universe. The reader is refer\\-red to \\cite{LaLu95} (1995) and to other contributions to this workshop for an introduction to cosmological topology and to \\cite{Lum98}~(1998) for an interesting historical introduction. A brief mathematical description of the relationship between a universal covering space $X,$ a compact 3-manifold $M$ and its fundamental polyhedron $P$ is provided in \\S\\ref{s-topol}. For a fuller introduction to three-dimensional geometry and topology, see \\cite{Thur82}~(1982, 1997). Three-dimensional topology detecting techniques are based on the required existence of multiple topological images of single physical objects. Techniques applicable to objects observed to successively larger scales are described in \\S\\ref{s-crystal} (``cosmic crystallography'', \\cite{LLL96}~1996), \\S\\ref{s-Xclus} (``brightest X-ray clusters'', \\cite{RE97}~1997) and \\S\\ref{s-isomet} (``local isometry detection'', \\cite{Rouk96}~1996). A Friedmann-Lema\\^{\\i}tre-Robertson-Walker metric (implying constant curvature of any spatial hypersurface) is assumed throughout this paper. Comoving coordinates are used, i.e., positions of objects observed in our past time cone are projected to the (3-D) space-like hypersurface at the present epoch, $t=t_0$. Spectroscopic redshifts, denoted $z,$ are used to obtain radial distance estimates (termed ``proper distances'' by \\cite{Wein72}~1972, eq.14.2.21\\footnote{The ``proper distance'' should not be confused with the quantity that \\protect\\cite{Wein72}~(1972, p.485) calls ``proper motion distance'' and that \\cite{Peeb93}~(1993, p.321, eq.13.36) calls ``angular size distance''. The proper distance and proper motion distance are identical for zero curvature, but not otherwise.}) in the $t=t_0$ space-like hypersurface. ", + "conclusions": "Several methods have been developed in the last few years to either detect or constrain the topology of the spatial part of the Universe. The relative efficiency (in terms of fundamental polyhedron crossings) of 3-D to 2-D methods depends moderately on the precise values of the metric parameters $\\Omega_0, \\lambda_0.$ Objects seen to about $z\\sim3$ would cross half the horizon distance for any presently accepted metric parameters, and in a cosmological constant dominated universe, objects seen to $z\\sim 0\\.1 -0\\.5$ would cover many fewer copies of the fundamental polyhedron than the CMB. Initial applications of 3-D methods to existing observational catalogues or individual observations indicate several (weak) candidates for the 3-manifold in which we live. (Or more precisely, for some of the generators of the 3-manifold.) These candidates are falsifiable with moderate observational investment in telescope time. Moreover, further development of the local isometry search method is presently possible for application to existing observational quasar catalogues. Looking to the future, new catalogues of objects made over the next few years, in particular all-sky surveys of quasars, will possibly allow the topology of the Universe to be detected to a high significance by the local isometry search method. Alternatively, the ``circles method'' of \\cite{Corn98b}~(1998b) applied to the observations by either MAP or Planck (planned CMB satellites) is likely to either reveal or constrain the topology of the Universe. Within a decade, we should know whether or not the topology of the Universe is detectable, and if so what it is. This research has been partially supported by the Polish Council for Scientific Research Grant KBN 2 P03D 008 13." + }, + "9802/astro-ph9802227_arXiv.txt": { + "abstract": "We carry out three-dimensional magnetohydrodynamical simulations of the magnetorotational (Balbus-Hawley) instability in weakly-ionized plasmas. We adopt a formulation in which the ions and neutrals each are treated as separate fluids coupled only through a collisional drag term. Ionization and recombination processes are not considered. The linear stability of the ion-neutral system has been previously considered by Blaes \\& Balbus (1994). Here we extend their results into the nonlinear regime by computing the evolution of Keplerian angular momentum distribution in the local shearing box approximation. We find significant turbulence and angular momentum transport when the collisional frequency is on order 100 times the orbital frequency $\\Omega$. At higher collision rates, the two-fluid system studied here behaves much like the fully ionized systems studied previously. At lower collision rates the evolution of the instability is determined primarily by the properties of the ions, with the neutrals acting as a sink for the turbulence. Since in this regime saturation occurs when the magnetic field is superthermal with respect to the ion pressure, we find the amplitude of the magnetic energy and the corresponding angular momentum transport rate is proportional to the ion density. Our calculations show the ions and neutrals are essentially decoupled when the collision frequency is less than $0.01\\Omega$; in this case the ion fluid behaves as in the single fluid simulations and the neutrals remain quiescent. We find that purely toroidal initial magnetic field configurations are unstable to the magnetorotational instability across the range of coupling frequencies. ", + "introduction": "The key to understanding accretion disk dynamics lies with the angular momentum transport mechanism. Since the molecular viscosity of disks is very low, some form of ``anomalous viscosity'' must be present. Although the precise nature of this anomalous viscosity has long been elusive, the discovery that differentially rotating systems are magnetohydrodynamically (MHD) unstable (Balbus \\& Hawley 1991) has led to the conclusion that fully ionized disks must be MHD turbulent. Because this magnetorotational instability is caused by angular momentum transport, the resulting turbulence has precisely the right character to transport angular momentum outwards (Balbus, Hawley \\& Stone 1996, hereafter BHS96), as required for disks to accrete. Since the MHD instability plays such a fundamental role in disks, the question naturally arises, what is its behavior when the plasma is not fully ionized? Protostellar and protoplanetary disks as well as other molecular disks are the venues where this question is particularly significant. Although it is tempting to assume that in the absence of MHD turbulence, purely hydrodynamical turbulence will rise to the task of transporting angular momentum, this now seems highly unlikely. BHS96 demonstrated, through a combination of analysis and simulation, that differentially rotating systems are both linearly and nonlinearly locally stable to hydrodynamic perturbations so long as the standard Rayleigh criterion is satisfied. Even if initiated ``by hand,'' hydrodynamic turbulence is not self-sustaining. Outward transport of angular momentum through a net Reynolds stress requires a specific average correlation between the radial velocity and the angular momentum fluctuations in the turbulent flow. Of fundamental importance is the interaction of these velocity fluctuations with the background mean flow. In differentially rotating systems the source of free energy is the angular velocity gradient. Angular momentum fluctuations, on the other hand, act to reduce the background angular {\\it momentum} gradient, which has the opposite sign from the angular velocity gradient. A positive value of the Reynolds stress, required to tap into the free energy of the system, acts as a sink term for the evolution of the angular velocity fluctuations that make up the Reynolds stress itself. Thus the turbulence is not self-sustaining. The results of BHS96 further imply that enhanced angular momentum transport is not the necessary outcome of turbulence. Because transport requires a high degree of correlation between the radial and azimuthal velocity fluctuations, turbulence that does not have its origin in the mean differentially rotating flow is unlikely to be an efficient source of angular momentum transport. This point is emphasized by Stone \\& Balbus (1996), whose numerical simulation showed that while vertical convection can generate turbulence, the resulting net radial angular momentum transport was {\\it inward} at a low rate (so long as the equator was kept hot by the boundary conditions of the simulation). Similar results were found by Cabot (1996). Simulations by Ryu \\& Goodman (1994) of the action of a parametric tidal instability on a disk show the generation of turbulence, but without internal transport. These studies provide compelling evidence that the idea of purely hydrodynamic turbulence as an {\\it angular momentum transport mechanism} in protostellar disks should be discarded. If turbulence is to be driven by the differential rotation, it must be MHD turbulence. If hydrodynamic turbulence is present, it must be driven by some source other than the background shear and generally will not transport angular momentum. Such considerations are particularly important if turbulence is important in protoplanetary disks for chondrule and planetesimal formation (e.g., Cuzzi, Dobrovolskis \\& Hogan 1996). It now appears that the number of different mechanisms for transporting angular momentum in disks is very limited. Magnetic fields and spiral wave disturbances (nonaxisymmetric distortions) are the only viable general mechanisms. In the latter category, and within the context of protostellar disks, nonaxisymmetric gravitational instabilities may be important. Global wave mechanisms, however, are not normally associated with the generation of local turbulence. The dynamics of a disk controlled by such global waves will be quite different from a model based upon the usual local $\\alpha$ prescription for viscosity, which is itself based upon the ansatz of transport by local turbulent stresses. For protostellar disks we are forced to reckon with questions of the effectiveness of magnetic field coupling. There is, of course, nothing about this issue that is unique to the magnetorotational instability. If the magnetic field is not well coupled to the fluid, then all potential MHD processes in the disk will be affected. It is difficult to imagine, for example, a scenario in which a weak field in a disk is stablized by low ion-neutral coupling, yet remains involved in, say, a kinematic dynamo. If a protostellar disk or rotating molecular cloud is sufficiently well-coupled to a weak magnetic field for the field to be important in any way, then the magnetorotational instability is an important dynamical factor. Under what circumstances, then, will the instability operate and produce MHD turbulence in protoplanetary disks where the ionization fraction is quite low? The first step towards answering this question was taken by Blaes \\& Balbus (1994, hereafter BB), who performed an axisymmetric linear stability analysis for a vertical (and toroidal) field in a weakly ionized plasma for a number of limits. Their major finding is that the linear magnetorotational instability is present so long as the ion-neutral collision frequency exceeds the local epicyclic frequency. This condition will be satisfied even for very small ionization fractions. They also found that azimuthal fields can reduce the observed growth rates, although such fields do not eliminate the instability. While a linear analysis can indicate the presence or absence of the instability, its effectiveness as a transport mechanism must be determined by its nonlinear evolution. The first numerical study of the two-fluid magnetorotational instability was carried out by Mac Low et al. (1995). They assume ionization-recombination equilibrium and consider the low-ionization limit, neglecting the ion pressure and inertia. This reduces the problem to a single-fluid (neutral) plus a diffusion term in the induction equation (ambipolar diffusion limit). Using the ZEUS code they carried out a series of two-dimensional simulations of the vertical flux tube problem (as in Hawley \\& Balbus 1991) for various ion-neutral coupling strengths. Although these simulations did not follow the evolution much beyond the linear stage, the results were in agreement with the stability analysis of BB in the appropriate limit. Essentially, so long as there are unstable wavelengths available, and the coupling between ions and neutrals is sufficiently strong, the instability behaves much like the single fluid case. The instability ceases to operate when the ambipolar diffusion rate becomes comparable to the growth rate of the instability, i.e., the field diffusion time is $< \\Omega^{-1}$, where $\\Omega$ is the disk orbital frequency. In another study, Brandenburg et al. (1995) investigated the ambipolar diffusion limit in three-dimensional simulations of a local, vertically stratified disk. They considered a case where the ambipolar diffusion time was long compared to the orbital time. They found that in this limit (i.e., ambipolar diffusion sufficiently small) the instability remains effective, and continues to generate self-sustained turbulence that transports angular momentum outward, albeit at a rate slightly reduced from the fully-coupled case. In another simulation, the diffusion time was set comparable to $\\Omega^{-1}$ and the turbulence decayed. These first results, while important, are only a beginning. To date, all the numerical studies have considered the behavior of the partially ionized system in the limit where the inertia of the ions can be completely neglected, and where the ion density is everywhere a fixed power-law function of the neutral density. In this paper we will approach the problem using a genuine two-fluid, ion-neutral evolution to examine effects where the ions are free to move relative to the neutrals. We will investigate the transition from the well-coupled regime, through critical coupling where the collision frequency is comparable to the epicyclic frequency, and down to the fully uncoupled limit. This parameter study should more clearly define the physical conditions for which full MHD turbulence and accompanying angular momentum transport can be expected, and those for which the bulk of the system dynamics are essentially hydrodynamic. The full range of conditions is likely to be of importance somewhere within protostellar or protoplanetary disks. In some regions of these disk systems, the ionization fractions may be quite small, but in other regions, such as near the forming protostar, the temperatures will be high, and the gas will be nearly fully ionized. In between, there will be a transition region. Determining the size of this transition region, and delineating its properties, will depend on obtaining a better understanding of the nonlinear behavior of the ion-neutral system in various regimes. The plan of the paper is as follows. In \\S2 we will consider the equations and the numerical techniques used to solve them. Although the collision term is handled semi-implicitly, the use of explicit finite-differencing for the remainder of the system makes the code Courant-limited. As a test problem, we compare numerical growth rates with analytic values from BB. In \\S3 we present the results of an extensive ensemble of simulations, covering a range of ionization fractions and coupling frequencies. Because of the Courant limit, this study is limited to relatively large ionization fractions $f \\equiv \\rho_i/(\\rho_i + \\rho_n)$. In the linear limit, however, the growth rates for small ionization fractions are relatively unaffected by decreasing $f$, if the ratio of the coupling frequency to orbital frequency is held constant. Thus, we expect that physical insights gained by this study should extend even into the small $f$ regime. The implications of the simulations will be summarized and discussed in \\S4. ", + "conclusions": "" + }, + "9802/astro-ph9802011_arXiv.txt": { + "abstract": "We report the measurement of the three-point correlation function (3PCF) of galaxies for the Las Campanas Redshift Survey (LCRS). We have not only measured the 3PCF in redshift space but also developed a method to measure the projected 3PCF which has simple relations to the real space 3PCF. Both quantities have been measured as a function of triangle size and shape with only a fractional uncertainty in each individual bin. Various tests derived from mock catalogs have been carried out to assure that the measurement is stable and that the errors are estimated reliably. Our results indicate that the 3PCFs both in redshift space and in real space have small but significant deviations from the well-known hierarchical form. The 3PCF in redshift space can be fitted by $\\Qsu=0.5\\cdot 10^{[0.2+0.1({s\\over s+1})^2]v^2}$ for $0.81\\mpc$ is noted. The real-space $\\Qru$ for $0.2 \\ls r_{12} \\ls 3\\mpc$ and $r_{31}\\ls 6\\mpc$ can be well described by {\\it half} the mean 3PCF predicted by a CDM model with $\\Omega_0 h=0.2$. The general dependence of the 3PCF on triangle shape and size is in qualitative agreement with the CDM cosmogonic models. Quantitatively the 3PCF of the models may depend on the biasing parameter and the shape of the power spectrum, in addition to other model parameters. Taking our result together with the constraints imposed by the two-point correlation function and the pairwise velocity dispersion of galaxies also obtained from the LCRS, we find that we have difficulties to produce a {\\it simple} model that meets all constraints perfectly. Among the CDM models considered, a flat model with $\\Omega = 0.2$ meets the 2PCF and PVD constraints, but gives higher values for the 3PCF than observed. This may indicate that more sophisticated bias models or a more sophisticated combination of model parameters must be considered. ", + "introduction": "The sample used for our analysis is the Las Campanas Redshift Survey (\\cite{setal96} 1996; hereafter LCRS). This is the largest redshift survey, which is now publicly available. Our sample consists of all galaxies with recession velocities between 10,000 and 45,000 $\\kms$ and with absolute magnitudes (in the LCRS hybrid R band) between $-18.0$ and $-23.0$. There are 19558 galaxies in this sample, of which 9480 are in the three north slices and the rest in the three south slices. The survey is a well-calibrated sample of galaxies, ideally suited for statistical studies of large-scale structure. All known systematic effects in the survey are well quantified and documented (Shectman et al. 1996; Lin et al. 1996), and so most can be corrected easily in statistical analyses. The only exception is the `fiber collision' limitation which prevents two galaxies in one $\\sim 1.5\\times 1.5\\, {\\rm deg}^2$ field from being observed when they are closer than $55''$ on the sky, because it is impossible to put fibers on both objects simultaneously. Here we will use extensively mock catalogs generated from N-body simulations to quantify this effect. The real-space 2PCF and the Pairwise Velocity Dispersion (PVD) have been determined for the LCRS by JMB98. The redshift-space 2PCF and power spectrum for this sample were presented by \\cite{tetal97} (1997) and \\cite{letal97} (1997) respectively. All these studies have shown that the LCRS is large enough to accurately measure these low order statistical quantities. In particular, JMB98 have carried out a detailed comparison between the observed 2PCF and PVD and the predictions of currently favoured CDM cosmogonies. They have used a large set of mock samples to adequately compare models and observations. The construction of mock catalogues from the simulations, i.e. photometric catalogues subject to the same selection effects as the real observations are a very important aspect of their analysis, because only in this way could the statistical significance of the results be asserted. Three spatially flat models have been considered in JMB98, with ($\\Omega_0$, $\\lambda_0$, $\\Gamma$,$\\sigma_8$)=(0.2,0.8,0.2,1.), (0.3,0.7,0.2,1.), and (1.0, 0.0, 0.5, 0.62), where $\\Omega_0$ is the density parameter, $\\lambda_0$ is the cosmological constant, $\\Gamma=\\Omega_0 h$ and $\\sigma_8$ are the shape parameter and normalization of the CDM power spectrum (\\cite{bbks86} 1986). All of the models give a steeper 2PCF, and a higher PVD on small scales than the data. Thus unless galaxies are biased with respect to the mass with a scale-dependent bias, all these models can be ruled out. Unfortunately physical models for a density or a (not so wanted, but perhaps unavoidable) velocity bias are not on firm grounds. Therefore in JMB98 a simple, but plausible phenomenological model for the bias has been suggested. To suppress the number of pairs in the DM distribution at small separations, it is assumed that the number of galaxies per unit dark matter mass $N/M$ is smaller in massive halos than in less massive ones. If a behaviour such as $N/M\\propto M_{cl}^\\alpha$ with $\\alpha=-0.08$ is used for clusters of mass $M_{cl}$, the predictions of some CDM models are consistent with the observational results. The best agreement was achieved for the flat $\\Omega_0=0.2$ model. We will use 10 mock catalogs of this model to test our statistical methods and quantify the `fiber collision' effect. Since this model has reproduced the LCRS 2PCF and PVD, we believe these mock catalogs are very suitable for this purpose. We shall also use these mock samples for model testing, as an example to illustrate the power of the three-point correlation function in discriminating between models which have similar two-point correlations. Since the model is a typical CDM model, we will generalize the discussion to other CDM models. ", + "conclusions": "The result is clear, and the conclusions are straightforward: We have succeeded in measuring the 3PCF from the LCRS. This is the first time that the three-point correlation function of galaxies has been measured accurately from a redshift survey. Both the 3PCF in redshift space and the projected 3PCF have been measured as a function of the triangle size and shape with only a fractional uncertainty in each individual bin. Various tests have been carried out to assure that the measurement is stable and that the errors are estimated reliably. Our results indicate that the 3PCFs both in redshift space and in real space have small but significant deviations from the well-known hierarchical form. The 3PCF in the redshift space can be fitted by $\\Qsu=0.5\\cdot 10^{[0.2+0.1({s\\over s+1})^2]v^2}$ for $0.8{E\\over 10^{15}\\ eV}{1\\over Z\\beta}, \\eqno(1)$$ where $L_{pc}$ is the size of the site in parsec, $Z$ is charge of the particle, and $\\beta$ is the speed of the scattering waves within the field of the site. The magnetic field needs to be large enough to confine the particles within their acceleration site and the size of the site must be sufficiently large for particles to gain sufficient energy before they escape. These simple requirements already rule out most astronomical objects in the universe. Figure \\ref{fig:accel_sites} shows objects which satisfy this dimensional requirement. Most of the galactic objects are excluded simply because they are too small and/or have magnetic fields that are too weak. Only a few extragalactic objects such as Active Galactic Nuclei (AGN) and radio galaxies remain as possible candidates. This fact is the basic reason many favor the extragalactic origin of EHECRs. In any actual acceleration site, energy loss mechanisms always compete with the gain of energy. With first order Fermi shock acceleration, the acceleration time is proportional to the mean free path for scattering in the shock wave, which itself is approximately inversely proportional to the magnetic field strength. Therefore, a certain magnitude of magnetic field is required, not only to confine the particles within the site, but also to accelerate the particles quickly. However, too strong a magnetic field also causes problems for particle acceleration, because it can cause protons to lose energy via synchrotron radiation. Other strong energy losses are caused by collisions with photons and/or matter at the acceleration site. A certain photon field is normally expected at the site as a result of synchrotron radiation by electrons or thermal radiation off the accretion disk. This leads to the additional requirement, that the site must have sufficiently low densities of radiation and matter so that cosmic ray nucleons are able to accelerate to $\\sim 10^{20}$ eV before losing significant energy. This raises more difficulties with candidate sites. For example, the core region of AGNs are ruled out because for this reason. The relativistic jets found in some classes of AGNs such as blazars may be able to produce EHECRs \\cite{ip91,mannheim95}, although models require optimistic fine-tuning of the acceleration efficiency and the Doppler boost factor of the relativistic jets. Rachen and Biermann have proposed hot spots of Fanaroff-Riley type II galaxies as EHECR sources \\cite{rachen93}. There seems to be enough acceleration power with not too-dense photons at the hot spots. However, the possibility is not excluded that collisions with UV photons in the spots discourage the acceleration of protons above $10^{20}$ eV. While relativistic jets of blazars and hot spots of FR II type radio galaxies are candidates sources of EHECRs via a conventional first order Fermi shock acceleration mechanism, it is not obvious that the acceleration efficiency is large enough to produce particles up to $10^{20}$ eV. These extragalactic models favor protons for the EHECR composition. Heavier nuclei like iron may break up into nucleons by photodisintegration during the shock acceleration, through collisions with UV photons at the site. \\bigskip The difficulties in acceleration can be avoided if EHECRs are direct products of processes which do not require acceleration. ``Top down'' scenarios have recently been proposed \\cite{bhattacharjee92,bhat95} involving relics of symmetry-breaking phase transitions in the early universe such as cosmic strings and magnetic monopoles, so called topological defects. If such defects exist, they may have produced EHE particles with energies up to the grand unified theory (GUT) scale (typically $\\sim 10^{25}$ eV) through the decay of the X-particles released in the collapse or decay of the defects. Because the hadron jets created at the decay of the X-particle are the main channels of particle production in this model, neutrinos and gamma rays, rather than protons and neutrons, are predominant. Any heavy nuclei like iron are completely ruled out in this model because the hadron jets create no nuclei. Propagation effects in the cosmic background radiation field (described later) would modify the emitting spectrum of each component, but one would still expect that gamma rays may be dominant at energies above $10^{20}$ eV, with details dependent on the strengths of the universal radio background and the extragalactic magnetic field, both of which are poorly known. The basic problem in this scenario, is that topological defects are exotic: the absolute intensity of defects is unknown. The observed intensities of cosmic rays and diffuse gamma rays can put constraints on the upper bound of defect intensity. So far, we have no experimental evidence, however, measurement of an excess of $\\gamma$ ray flux above $10^{20}$ eV and detection of EHE neutrinos above $10^{19}$ eV would be signatures of topological defects. It has been suggested that Gamma Ray Bursts (GRB), responsible for gamma rays up to the GeV range, may also be able to produce EHECRs. This would be a burst source and not a continuously emitting one. This would also result in a correlation between arrival times and energies of EHECRs. Unfortunately, the time scale might be much longer than any single experiment can afford to run and thus the correlation may be extremely difficult to detect. Detail on this idea is presented in \\cite{waxman95,vietri95}. \\subsection{Propagation of EHECRs in space} It is important to understand how EHECRs propagate from their sources to earth, since this puts constraints on possible sources and provides hints for the most effective way of searching for them. First, the galactic magnetic field of $\\sim\\mu$G can no longer confine cosmic ray protons with energies greater than $10^{19}$ eV in the galactic disk since the Larmor radius of a proton at that energy, $$ L_{kpc} \\simeq ({E\\over 10^{19}\\ eV})({B\\over 3\\ \\mu G})^{-1} \\times 3 kpc, \\eqno(2) $$ becomes greater than the thickness of our galactic disk. This means that any galactic protons can easily escape from our galaxy, provided that the galactic magnetic fields do not extend out into a possible Halo. This again favors the hypothesis of an energetic extragalactic component dominating galactic components in the EHECRs population. Secondly, when EHECRs are traveling through extragalactic space, their trajectories are not strongly bent by the extragalactic magnetic field and the arrival directions of such cosmic rays should point back to their emitters. Information on the extragalactic magnetic field strength is difficult to gather. We know only the Faraday rotation bound on the extragalactic magnetic field is given by \\cite{kronberg94} $$ B_{rms}\\sqrt{l_c} \\leq 10^{-9} G\\ Mpc^{1\\over 2} \\eqno(3) $$ where $l_c$ is the scale of the coherent magnetic field in Mpc, and the mean deflection angle can be written as $$ \\theta_{def} \\leq ({R\\over 10Mpc})^{1\\over 2} ({E\\over 3\\times 10^{19}\\ eV})^{-1} \\times 3.2^{\\circ} \\eqno (4) $$ for protons \\cite{waxman96}. Here $R$ is distance to the source. This opens a new window of astronomy, that of Charged Particle Astronomy. A typical deflection angle of $\\sim 3^{\\circ}$, which is comparable with the typical angular resolution of the present experiments, might be still too large in an actual search for sources, because there exist many astronomical objects in a $3^{\\circ}\\times 3^{\\circ}$ window even if the candidates are limited to AGNs and radio galaxies. The real situation is much better, however, since there is a limit on the distance over which an EHECR may travel. We can limit the search to relatively nearby sources, because EHECRs collide with cosmological backgrounds and lose energy during their propagation. This is the most important effect on the propagation of EHECRs. \\begin{figure} \\centerline{{\\epsfxsize7cm\\epsfbox{fig3.ps}}} \\caption{The attenuation length of cosmic rays as a function of energy. The solid curve shows the case for nucleons calculated by Yoshida and Teshima. The dashed curve shows the case for iron calculated by Puget, Stecker and Bredekamp. The bound given by redshift (adiabatic energy loss) are applicable to all primaries.} \\label{fig:att_length} \\end{figure} \\begin{figure} \\centerline{{\\epsfxsize6cm\\epsfbox{fig4.ps}}} \\caption{The expected spectral shapes for two assumptions concerning the distribution of sources in extragalactic space. $m$ is the evolution parameter of EHECR emission. A larger $m$ means more contributions from sources in distant (high redshift) regions. This extragalactic population might dominate in the highest energy region of the observed spectrum above $\\sim 10^{19}$ eV. The galactic magnetic field, which could suppress the the recovery below $10^{18.5}$ eV, is not included in the simulation.} \\label{fig:spec_proton_dif} \\end{figure} \\bigskip EHECR protons or neutrons interact with the microwave background photons through pair creation and photopion production. The threshold energy of photopion production, the main energy loss process, is $$ E \\simeq 7\\times 10^{19} ({E_{bb}\\over 2\\times 10^{-3}\\ eV})^{-1} (1+\\cos\\theta)^{-1}\\quad eV \\eqno(5) $$ where $E_{bb}$ is the energy of the microwave background photon, from a blackbody spectrum with a characteristic temperature of $2.7^{\\circ}$ K. The photon and the EHECR interact with a collision angle $\\theta$. Above the threshold energy, EHECRs rapidly lose energy. This may result in a cutoff in the energy spectrum. This cutoff is known as Greisen-Zatsepin-Kuzmin cutoff (GZK cutoff) \\cite{grei66,zats66}, and is the centerpiece of the EHECRs physics. A detection of this effect proves the extragalactic origin of EHECRs and limits the distance to possible sources to less than $\\simeq 100$ Mpc for particles above $10^{20}$ eV. The situation is explained in figure \\ref{fig:att_length} which shows the attenuation length of protons in extragalactic space. One finds that the attenuation length above the photopion production threshold is contracted by rapid energy losses. The attenuation length for protons with energies higher than $7\\times 10^{19}$ eV (the threshold energy of photopion production) is shorter than 500 Mpc \\cite{yosh93,prot96}. Any sources contributing to the bulk of EHECRs above this energy should be within 500 Mpc of earth. The higher the energy, the shorter the upper bound on the distance. A $3\\times 10^{20}$ eV proton would require sources within only $\\sim 50$ Mpc. Thus, nearby sources should make the dominant contributions at the high energy end of the spectrum. This effect provides an important feature on the resulting energy spectrum {\\it shape}. Because the microwave background during cosmological evolution is a function of redshift, the spectral shape of EHECRs also depends on the redshift, as well as the source distribution in space. Figure \\ref{fig:spec_proton_dif} shows the expected spectral shapes if many sources are isotropically distributed in the universe \\cite{yosh93}. The parameter $m$ describes the cosmological evolution of cosmic ray emission. Therefore, it controls the relative contributions of sources at different distances. The spectral shape changes with the parameter $m$, however, the cutoff energy remains near $5\\times 10^{19}$ eV. The dominant contribution of nearby sources at the high energy end make the spectral shape above $10^{19}$ eV less sensitive to cosmological effects. The shape around the GZK cutoff is universal while most of the cosmological signatures are found in the $10^{17}-10^{18}$ eV region where another cosmic ray population may dominate. A search for a cutoff at around $5\\times 10^{19}$ eV is indeed a robust method for obtaining evidence of the extragalactic origin regardless of details in the model. A similar situation exists for primary nuclei, like carbon or iron. This time, photodisintegration is the limiting factor rather than photopion production. As a result, there is an even more rapid energy loss during propagation as shown in figure \\ref{fig:att_length} \\cite{puge76}. Since heavy nuclei break down quickly during propagation, an EHECR composition favoring protons and neutrons is likely. EHE nuclei will be reduced both at the acceleration sites and over the propagation volume. \\begin{figure} \\centerline{{\\epsfxsize9cm\\epsfbox{fig5.ps}}} \\caption{The attenuation length of photons traveling in extragalactic space.} \\label{fig:photon_att} \\end{figure} \\bigskip We should also consider the possibility of EHE cosmic rays being photons. The most conventional mechanism for the production of EHE photons is the decay of neutral pions produced by a collision between an EHE cosmic ray nucleon and a background photon during propagation. The exotic ``top down'' scenario involving topological defects predicts a more predominant initial photon flux \\cite{sigl96}. These EHE photons/electrons initiate electro-magnetic cascades on a low energy radiation field such as the microwave background. The attenuation length of photons in the radiation field of extragalactic space is shown in figure \\ref{fig:photon_att} \\cite{lee96}. EHE photons interact with microwave/radio background photons via pair creation and double pair creation. Electrons produced in this process transfer most of their energy to a background photon via inverse Compton scattering or sometimes via triplet pair production ($e\\gamma_b \\to e e^+e^-$). Since the EHE $\\gamma$ ray attenuation length does not decrease with energy (as is the case for protons), there is no cutoff feature in the spectrum \\cite{sigl96,lee96}. This leads to the prediction of a dominant gamma ray flux at energies above the GZK cutoff. It should be pointed out that the $\\gamma$ ray flux depends on two poorly known parameters: the extragalactic magnetic field and the universal radio background. A strong radio field reduces the mean free path for pair creation, and synchrotron radiation cools the electron pairs out of the EHE range. The conventional shock acceleration models always predict very low gamma ray fluxes \\cite{yosh93} while the ``top down'' models provide a possibility of $\\gamma$ ray dominance around $10^{20}$ eV\\cite{sigl96}. \\begin{figure} \\centerline{{\\epsfxsize7cm\\epsfbox{fig6.ps}}} \\caption{The horizon of the universe for EHE neutrinos as a function of present-day neutrino energies. Two cases are shown for neutrino primary energies at emission of $10^{15}$ GeV (thin lines) and $10^{16}$ GeV (thick lines), both of which could be reasonable in the top down model of EHECR production involving topological defects.} \\label{fig:neut_horizon} \\end{figure} \\bigskip Finally, we discuss the case of neutrinos. EHE neutrinos can certainly be created through the decay of charged pions produced by collisions between EHE cosmic ray nucleons and microwave background photons \\cite{yosh93,hill83,yoshida97}. These secondary neutrinos are good probes of EHE particle emission activities at early epochs of the universe, since their flux strongly depends on the evolution parameters. The detection of these neutrinos is unfortunately a remote possibility since the most optimistic flux is only comparable to that of observed EHE cosmic rays. Nevertheless, a search for neutrinos in the EHE range (above $\\sim 10^{19}$ eV) is a meaningful test of the topological defect hypothesis since that model predicts a much higher neutrino flux \\cite{yoshida97,sigl97}. Because the maximum energy of neutrinos reaches the GUT scale ($\\sim 10^{25}$ eV) and their emission at superhigh redshift epochs ($z\\sim 500$) are the main contributions in this scenario, collisions of EHE neutrinos with low energy cosmological relic neutrinos are not negligible \\cite{roulet93,yoshida94}. Figure \\ref{fig:neut_horizon} shows the horizon of the universe for EHE neutrinos, the maximum redshift to which EHE neutrinos are not attenuated in their propagation \\cite{yoshida94}. The dotted lines correspond to the upperbound of the horizon when one considers the redshift energy loss only. It is shown that interactions with the relic neutrinos, which contract the horizon, are a key effect in the ``top down'' scenario because the emissions of EHE neutrinos at high redshift epochs ($z>100$) are predominant due to the higher rate of annihilation of topological defects. It should be noted that some EHE neutrinos that initiate neutrino cascading on the cosmological neutrino background field will further enhance the EHE neutrino flux at earth, and the planned future experiments may be able to detect a few of them \\cite{yoshida97}. These interactions also play an important role in the recently proposed mechanism to generate observable particles above the GZK cutoff \\cite{weiler97} by collisions of EHE cosmic ray neutrinos with possibly clustered massive neutrinos in our galaxy. ", + "conclusions": "The EHE cosmic ray energy spectrum steepens in the energy region between $10^{17.6}$ and $10^{18}$eV (the second ``knee'', where the spectral slope changes from -3.0 to -3.3) and flattens between $10^{18.5}$ and $10^{19}$ eV (the so called ``ankle'', where the spectral slope changes from -3.3 to approximately -2.7). The straightforward, and less model-dependent, interpretation is a two component scenario: a high energy extragalactic component dominates over a steeper galactic component above the ankle. The many experimental results now available are supportive of this picture, including: a possible signature of the GZK cutoff obtained by the Fly's Eye, AGASA and Yakutsk; an indication of the chemical composition getting lighter at high energies from the Fly's Eye and Yakutsk groups; no enhancement of the arrival direction distribution associated with the galactic plane; a possible correlation with the supergalactic plane found in a combined data set mainly consisting of Haverah Park data; and three event clusters above $4\\times 10^{19}$ eV observed by AGASA arriving from directions well away from the galactic plane. Although our two component picture seems to make sense when we put these results {\\it together}, the conclusion is far from solid. The significance of the GZK cutoff is muddied by the super-high energy events well beyond the cutoff, thereby providing complications to the simple picture of the GZK mechanism. The interpretation of chemical composition measurements has some model dependence as cautioned by the AGASA results. The statistical significance of the event clusters only allows us to suggest possible ``hints'' of something exciting. What encourages us about the two component scenario is the fact that different analysis from different experiments seem to be reasonably consistent under our scenario. The next step is to make all these results robust by accumulating more data with good resolution. For example, a fine measurement of the GZK cutoff would clarify the extragalactic hypothesis. A clear measurement of the mass composition above the ``ankle'' would also be very helpful in confirming or rejecting our current picture. A detection of a EHE photon or neutrino component would bring us a new understanding of the universe. During the next decade, we will see the study of EHE cosmic rays continue to provide a laboratory for non-accelerator particle physics, and we look forward to it establishing a new astronomy. \\vspace{0.5cm} \\centerline{\\Large{\\bf Acknowledgments}} The authors are grateful to B. Dawson, P. Sommers, P. Sokolsky, M. Teshima, J. N. Matthews, and G. B. Yodh for helpful suggestions and advice. They also wish to thank M. Nagano for allowing us to use the preliminary analysis by one of the authors (S.Y.) based on the published data from the AGASA experiment." + }, + "9802/gr-qc9802038_arXiv.txt": { + "abstract": "We discuss a dramatic difference between the description of the quantum creation of an open universe using the Hartle-Hawking wave function and the tunneling wave function. Recently Hawking and Turok have found that the Hartle-Hawking wave function leads to a universe with $\\Omega = 0.01$, which is much smaller that the observed value of $\\Omega$. Galaxies in such a universe would be $10^{10^8}$ light years away from each other, so the universe would be practically structureless. We argue that the Hartle-Hawking wave function does not describe the probability of creation of the universe. If one uses the tunneling wave function for the description of creation of the universe, then in most inflationary models the universe should have $\\Omega = 1$, which agrees with the standard expectation that inflation makes the universe flat. The same result can be obtained in the theory of a self-reproducing inflationary universe, independently of the issue of initial conditions. However, there exist some models where $\\Omega$ may take any value, from $\\Omega > 1$ to $\\Omega \\ll 1$. ", + "introduction": "\\subsection{Why do we need quantum cosmology?} The investigation of the wave function of the universe goes back to the fundamental papers by Wheeler and DeWitt \\cite{DeWitt}. However, for a long time it seemed almost meaningless to apply the notion of the wave function to the universe itself, since the universe is not a microscopic object. Only with the development of inflationary cosmology it became clear that the whole universe could appear from a tiny part of space as small as the Planck length $M_p^{-1}$ (at least in the chaotic inflation scenario \\cite{Chaotic}). Such a tiny region of space can appear as a result of quantum fluctuations of metric, which should be studied in the context of quantum cosmology. Later it was found that the global structure of the universe in the chaotic inflation scenario is determined not by classical physics, but by quantum processes \\cite{book}. Unfortunately, quantum cosmology is not a well developed science. This theory is based on the Wheeler-DeWitt equation, which is the Schr\\\"{o}dinger equation for the wave function of the universe. This equation has many solutions, and at the present time the best method to specify preferable solutions of this equation, as well as to interpret them, is based on the Euclidean approach to quantum gravity. This method is very powerful, but some of its applications are not well justified. In some cases this method may give incorrect answers, but rather paradoxically sometimes these answers appear to be correct when applied to some other questions. Therefore it becomes necessary not only to solve the problem in the Euclidean approach, but also to check, using one's best judgement, whether the solution is related to the original problem or to something else. An alternative approach is based on the use of stochastic methods in inflationary cosmology \\cite{Star,Gonch,Star2,book,LLM}. These methods allow one to understand such effects as the creation of inflationary density perturbations, the theory of tunneling, and even the theory of self-reproduction of inflationary universe. Both Euclidean approach and stochastic approach to inflation have their limitations, and it is important to understand them. \\subsection{Hawking-Moss tunneling} Before discussing quantum creation of the universe, let us pause a little and study the problem of tunneling between two local minima of the effective potential $V(\\phi)$ in inflationary cosmology. As we will see, this subject is closely related to the issue of quantum creation of the universe. Consider a theory with an effective potential $V(\\phi)$ which has a local minimum at $\\phi_0$, a global minimum at $\\phi_*$ and a barrier separating these two minima, with the top of the barrier positioned at $\\phi = \\phi_1$. One of the first works on inflationary cosmology was the paper by Hawking and Moss \\cite{HM} where they studied a possibility of tunneling from $\\phi_0$ to $\\phi_*$ in the new inflationary universe scenario. They have written equations of motion for the scalar field in an Euclidean space with the metric \\begin{equation}\\label{metric} ds^2 =d\\tau^2 +a^2(\\tau)(d \\psi^2+{\\rm sin}^2 \\psi d \\Omega_2^2) \\ . \\end{equation} The field $\\phi$ and the radius $a$ obey the field equations \\begin{equation}\\label{equations} \\phi''+3{a'\\over a}\\phi'=V_{,\\phi},~~~~~ a''= -{8\\pi G\\over 3} a ( \\phi'^2 +V) \\ , \\end{equation} where primes denote derivatives with respect to $\\tau$. If the potential has an extremum at some particular value of the field $\\phi$, then the equation for the field $\\phi$ is solved trivially by the field staying at this extremum. Then the equation for $a(\\tau)$ has a simple solution $a(\\tau)=H^{-1} {\\sin} (H\\tau)$, with $H^2= 8 \\pi G V(\\phi)/3 = 8 \\pi V(\\phi)/3M_p^2$. This solution describes a sphere $S^4$, the Euclidean version of de Sitter space. In this description $\\tau$ plays the role of Euclidean time, and $a(\\tau)$ the role of the scale factor. One can try to interpret one half of this sphere as an instanton. The action on this instanton is negative, \\begin{equation}\\label{action} S = \\int d^4 x \\sqrt{-g}\\left(-{RM_p^2\\over{16\\pi }}+V(\\phi)\\right) = - {3M_p^4\\over 16 V(\\phi)} \\ . \\end{equation} It was argued in \\cite{HM} that the probability of tunneling from $\\phi_0$ to the true vacuum $\\phi_*$ is given by \\begin{equation}\\label{HMinst} P \\sim \\exp\\left({3M_p^4\\over 8V(\\phi_1)}\\right) \\ \\exp\\left(-{3M_p^4\\over 8 V(\\phi_0)}\\right) {}~. \\end{equation} The probability of tunneling, as usual, is suppressed by $e^{-2S}$ (or by $e^{-S}$ if by $S$ we mean the result of integration over the whole sphere, $-{3M_p^4\\over 8V(\\phi)}$). This is the standard result of the Euclidean theory of tunneling. Everything else about this result was rather mysterious. First of all, instantons typically interpolate between the initial vacuum state and the final state. Here, however, the scalar field on the instanton solution was exactly constant. So why do we think that they describe tunneling from $\\phi_0$ if $\\phi_0$ never appears in the instanton solution? \\begin{figure}[Fig0111] \\hskip 1.5cm \\leavevmode\\epsfysize=7cm \\epsfbox{HawkMoss.eps} \\ \\caption[Fig1]{\\label{HawkMoss} A possible interpretation of the Hawking-Moss tunneling from $\\phi_0$ to $\\phi_1$.} \\end{figure} A possible answer to this question can be given as follows. One can choose the coordinate system where inflationary universe looks as a closed de Sitter space near the point of a maximal contraction, where its size becomes $H^{-1}(\\phi_0)$, see region 1 in Fig. \\ref{HawkMoss}. Classically, such a universe at that moment begins expanding with the same value of the Hubble constant as before. However, since the total size of the universe at that moment is finite, it may also jump quantum mechanically to a state with a different value of the field $\\phi$ corresponding to a different extremum of the effective potential. One can represent this process by gluing two de Sitter instantons corresponding to two different values of the scalar field $\\phi$ ($\\phi_0$ in the region 2, and $\\phi_1$ in the region 3 in Fig. \\ref{HawkMoss}), and by making analytical continuation to the Lorentzian regions 1 and 4. This seems to be a plausible interpretation of the Hawking-Moss tunneling (see also \\cite{CB}). But it certainly does not answer all questions. What will happen if we have several different local minima and maxima of $V(\\phi)$? Why does the tunneling go to the top of the effective potential rather than to the absolute minimum of the effective potential, or to some other local maximum? Finally, if the instanton describes an exactly homogeneous scalar field $\\phi$, does it mean that the tunneling must simultaneously occur everywhere in an exponentially large inflationary universe? This does not seem plausible, but what else should we think about, if the field $\\phi$ on the instanton solution is constant? And indeed, originally it was assumed that the tunneling described by this instanton must occur simultaneously in the whole universe. Then, in the second paper on this subject, Hawking and Moss said that their results were widely misunderstood, and that this instanton describes tunneling which is homogeneous only on the scale of horizon $\\sim H^{-1}$ \\cite{HM2}. But how is it possible to describe inhomogeneous tunneling by a homogeneous instanton? A part of the answer was given in Ref. \\cite{Gonch}. We have found that if one deforms a little the Hawking-Moss instanton to make the field $\\phi$ match $\\phi_0$ in some small region of the sphere, we will, strictly speaking, not get a solution, but the action on such a configuration can be made almost exactly coinciding with the Hawking-Moss action. Then such configurations can play the same role as instantons \\cite{Nucl}. A full understanding of this issue was reached only after the development of the stochastic approach to inflation \\cite{Star,Gonch,Star2,book}. We will return to this question later. \\subsection{Creation of the universe from nothing} Now we will discuss the problem of the universe creation. According to classical cosmology, the universe appeared from the singularity in a state of infinite density. Of course, when the density was greater than the Planck density $M_p^4$ one could not trust the classical Einstein equations, but in many cases there is no demonstrated need to study the universe creation using the methods of quantum theory. For example, in the simplest versions of the chaotic inflation scenario \\cite{Chaotic}, the process of inflation, at the classical level, could begin directly in the initial singularity. However, in certain models, such as the Starobinsky model \\cite{b14} or the new inflationary universe scenario \\cite{New}, inflation cannot start in a state of infinite density. In such cases one may speculate about the possibility that inflationary universe appears due to quantum tunneling ``from nothing.'' The first idea how one can describe creation of an inflationary universe ``from nothing'' was given in 1981 by Zeldovich \\cite{Zeld} in application to the Starobinsky model \\cite{b14}. His idea was qualitatively correct, but he did not propose any quantitative description of this process. A very important step in this direction was made in 1982 by Vilenkin \\cite{NothVil}. He suggested to calculate the Euclidean action on de Sitter space with the energy density $V(\\phi)$, which coincides with the Hawking-Moss instanton with the action $S = -{3M_p^4\\over 16 V(\\phi)}$. However, as we have seen, this instanton by itself does not tell us where the tunneling comes from. Vilenkin suggested to interpret this instanton as the tunneling trajectory describing creation of the universe with the scale factor $a = H^{-1} = \\sqrt{3 M_{\\rm P}^2\\over 8\\pi V}$ from the state with $a = 0$. This would imply that the probability of quantum creation of the universe is given by \\begin{equation} {P} \\propto \\exp (-2S) = \\exp \\left({3 M_p^4 \\over 8 V(\\phi)}\\right). \\label{Vil1} \\end{equation} A year later this result received strong support when Hartle and Hawking reproduced it by a different though closely related method \\cite{HH}. They argued that the wave function of the ``ground state'' of the universe with a scale factor $a$ filled with a scalar field $\\phi$ in the semiclassical approximation is given by \\begin{equation}\\label{E31} \\Psi_0(a,\\phi)\\sim \\exp\\left(-S(a,\\phi)\\right)\\ . \\end{equation} Here $S(a,\\phi)$ is the Euclidean action corresponding to the Euclidean solutions of the Lagrange equation for $a(\\tau)$ and $\\phi(\\tau)$ with the boundary conditions $a(0)=a, \\phi(0)=\\phi$. The reason for choosing this particular wave function was explained as follows. Let us consider the Green's function of a particle which moves from the point $(0,t')$ to the point ${\\bf x},t$: \\begin{eqnarray}\\label{E32} <{\\bf x},t|0, t'> &=& \\sum_n \\Psi_n ({\\bf x})\\Psi_n(0) \\exp\\left(iE_{n}(t-t')\\right) \\nonumber \\\\ &=& \\int d{\\bf x}(t) \\exp\\left(iS({\\bf x}(t))\\right)\\ , \\end{eqnarray} where $\\Psi_n$ is a complete set of energy eigenstates corresponding to the energies $E_n\\geq 0$. To obtain an expression for the ground-state wave function $\\Psi_0({\\bf x})$, one should make a rotation $t \\rightarrow -i\\tau$ and take the limit as $\\tau \\rightarrow -\\infty$\\@. In the summation (\\ref{E32}) only the term $n=0$ with the lowest eigenvalue $E_0 = 0$ survives, and the integral transforms into $\\int dx(\\tau)\\exp(-S({\\bf x}(\\tau)))$. This yields, in the semiclassical approximation, \\begin{equation}\\label{E31aa} \\Psi_0(x)\\sim \\exp\\left(-S({\\bf x})\\right)\\ , \\end{equation} where the action is taken on the classical trajectory bringing the particle to the point ${\\bf x}$. Hartle and Hawking have argued that the generalization of this result to the case of interest would yield (\\ref{E31}). The method described above is very powerful. For example, it provides the simplest way to find the wave function of the ground state of the harmonic oscillator in quantum mechanics. However, this wave function simply describes the probability of deviations of the harmonic oscillator from its equilibrium. It certainly does not describe quantum creation of a harmonic oscillator. Similarly, if one applies this method to the hydrogen atom, one can obtain the wave function of an electron in the state with the lowest energy. Again, this result has no relation to the probability of creation of an electron from nothing. The gravitational action involved in (\\ref{E31}) is the same action as before, corresponding to one half of the Euclidean section $S_4$ of de Sitter space with $a(\\tau) = H^{-1}(\\phi)\\cos H\\tau$ ($0\\leq\\tau\\leq H^{-1}$). One can represent it in the following form: \\begin{eqnarray}\\label{E33} S(a, \\phi) &=& - \\frac{3\\pi M_p^2}{4} \\int d\\eta\\Bigl[\\Bigl(\\frac{da}{d\\eta}\\Bigr)^2 - a^2 + \\frac{8\\pi V}{3M_p^2}a^4\\Bigr] \\nonumber \\\\ &=& - \\frac{3M_p^4}{16 V(\\phi)}\\ . \\end{eqnarray} Here $\\eta$ is the conformal time, $\\eta = \\int {d\\tau\\over a(\\tau)}$. Therefore, according to \\cite{HH}, \\begin{equation}\\label{E34} \\Psi_0(a,\\phi)\\sim \\exp{\\Bigl(-S(a,\\phi)\\Bigr)} \\sim \\exp\\left(\\frac{3 M_p^4 }{16V(\\phi)}\\right) . \\end{equation} By taking a square of this wave function one again obtains eq. (\\ref{Vil1}). The corresponding expression has a very sharp maximum as $V(\\phi) \\rightarrow 0$. This could suggest that the probability of finding the universe in a state with a large field $\\phi$ and having a long stage of inflation should be strongly suppressed. But is it a correct interpretation of the Hartle-Hawking wave function? Just like in the examples with the harmonic oscillator and the hydrogen atom mentioned above, nothing in the `derivation' of the Hartle-Hawking wave function tells that it describes creation of the universe from nothing. The simplest way to interpret the Hartle-Hawking wave function in application to de Sitter space is as follows. At the classical level, de Sitter space has a definite speed of expansion, definite size of its throat $H^{-1}$, etc. At the quantum level, de Sitter ``trajectory'' becomes wider because of quantum fluctuations. The Hartle-Hawking wave function of de Sitter space describes the probability of deviations of metric of de Sitter space from its classical expectation value, which may occur due to the process shown in Fig. \\ref{HawkMoss}. This is very much different from the probability of spontaneous creation of the universe. In fact, Eq. (\\ref{Vil1}) from the very beginning did not seem to apply to the probability of creation of the universe. The total energy of matter in a closed de Sitter space with $a(t) = H^{-1}\\cosh Ht$ is greater than its minimal volume $\\sim H^{-3}$ multiplied by $V(\\phi)$, which gives the total energy of the universe $E {\\ \\lower-1.2pt\\vbox{\\hbox{\\rlap{$>$}\\lower5pt\\vbox{\\hbox{$\\sim$}}}}\\ } M_p^3/\\sqrt V$. Thus the minimal value of the total energy of matter contained in a closed de Sitter universe {\\it grows} when $V$ decreases. For example, in order to create the universe at the Planck density $V \\sim M_p^4$ one needs no more than the Planckian energy $M_p \\sim 10^{-5}$ g. For the universe to appear at the GUT energy density $V \\sim M_X^4$ one needs to create from nothing the universe with the total energy of matter of the order of $M_{\\rm Schwarzenegger} \\sim 10^2$ kg, which is obviously much more difficult. Meanwhile, if one makes an attempt to use the Hartle-Hawking wave function for the description of the creation of the universe (which, as we believe, does not follow from its derivation), then eq. (\\ref{Vil1}) suggests that it should be much easier to create a huge universe with enormously large total mass rather than a small universe with Planckian mass. This seems very suspicious. From uncertainty relations one can expect that the probability of a process of universe formation is not exponentially suppressed if it occurs within a time $\\Delta t < E^{-1}$. This is quite possible if the effective potential is of the order of $M_p^4$ and $E \\sim M_p^3/\\sqrt V \\sim M_p$. In such a case one may envisage the process of quantum creation of a universe of mass $M_p$ within the Planck time $M_p^{-1}$. However, the universe of mass $E \\gg M_p$ (which is the case for $V \\ll M_p^4$) can be created only if the corresponding process lasts much shorter than the Planck time $M_p^{-1}$, which is hardly possible. Another way to look at it is to calculate the total entropy ${\\rm \\bf S}$ of de Sitter space at the moment of its creation. It is equal to one quarter of the horizon area of de Sitter space (in Planck units), which gives ${\\rm \\bf S} = {3 M_p^4\\over 8 V(\\phi)}$. (Note its relation to the Euclidean action on the full de Sitter sphere ${ S} = -{3 M_p^4\\over 8 V(\\phi)}$.) It seems natural to expect that the probability of emergence of a complicated object of large entropy must be suppressed by a factor of $\\exp({-\\rm \\bf S}) = \\exp(-{3 M_p^4\\over 8 V(\\phi)})$, which again brings us to the equation (\\ref{TUNN}), see \\cite{KLB}. Meanwhile the use of the Hartle-Hawking wave function for the description of creation of the universe would indicate that it is much more probable to create a very large universe with a huge entropy rather than a small universe with entropy $O(1)$. To avoid misunderstandings, one should note, that the probability of fluctuations in a thermodynamical system is always {\\it suppressed} by the factor $e^{\\rm \\bf \\Delta S}$, where ${\\rm \\bf \\Delta S}$ is the change of entropy between two different states of the system \\cite{HTnew}. As we will see, this is exactly what happens during the tunneling between two different states of de Sitter space with two different values of $V(\\phi)$. This is in perfect agreement with the prediction of the Hartle-Hawking wave function if one applies it not to the creation of the universe but to the probability of its change. However, now we are not talking about the probability of change of the state of the system, but about a possibility of creation of the whole system together with a lot of information stored in it from nothing. We are not going to insist that this process is possible. In fact in chaotic inflation scenario this assumption is not necessary because the universe formally can inflate even in a state with indefinitely large density, so there is no need for any tunneling to take place. However, if creation from nothing is possible at all, then the tunneling wave function suggests that this process should be as unintrusive as possible, whereas the Hartle-Hawking approach implies that the greater the change, the easier it occurs. I leave it for the reader to decide whether this looks plausible. One may wonder why the Hartle-Hawking wave function leads to rather counterintuitive predictions when applied to the probability of creation of the universe? There is one obvious place where the derivation (or interpretation) of eq. (\\ref{Vil1}) could go wrong. The effective Lagrangian of the scale factor $a$ in (\\ref{E33}) has a wrong overall sign. Solutions of the Lagrange equations do not know anything about the sign of the Lagrangian, so we may simply change the sign before studying the tunneling. Only after switching the sign of the Lagrangian of the scale factor in (\\ref{E33}) and representing the theory in a conventional form can we consider tunneling of the scale factor. But after changing the sign of the action, one obtains a different expression for the probability of quantum creation of the universe: \\begin{equation} {P} \\propto \\exp (-2|S|) = \\exp \\left(-{3 M_p^4\\over 8 V(\\phi)}\\right). \\label{E366} \\end{equation} This equation predicts that a typical initial value of the field $\\phi$ is given by $V(\\phi)\\sim M_p^4$ (if one does not speculate about the possibility that $V(\\phi) \\gg M_p^4)$, which leads to a very long stage of inflation. Originally I obtained this result by the method described above. However, because of the ambiguity of the notion of tunneling from the state $a = 0$, one may try to look at the same subject from a different perspective, and reexamine the derivation of the Hartle-Hawking wave function. In this case the problem of the wrong sign of the Lagrangian appears again, though in a somewhat different form. Indeed, the total energy of a closed universe is zero, being a sum of the positive energy of matter and the negative energy of the scale factor $a$. Thus, the energy $E_n$ of the scale factor is negative. If one makes the same Euclidean rotation as in Eq. (\\ref{E32}), the contributions of all states with $n >1$ will be greater than the contribution of the state with the lowest absolute value of energy, so such a rotation would not allow one to extract the wave function $\\Psi_0$ as we did before. This is a simple mathematical fact, which means that the main argument used in \\cite{HH} to justify their prescription of quantization of the scale factor fails. In order to suppress terms with large negative $E_n$ and to obtain $\\Psi_0$ from (\\ref{E32}) one should rotate $t$ not to $-i\\tau$, but to $+i\\tau$. This gives \\cite{Creation} \\begin{equation}\\label{E35a} \\Psi_0(a,\\phi) \\sim \\exp\\Bigl(-|S(a,\\phi)|\\Bigr) \\sim\\exp \\left(- \\frac{3 M_p^4}{16V(\\phi)}\\right) , \\end{equation} and \\begin{equation}\\label{E36} P(\\phi) \\sim|\\Psi_0(a,\\phi)|^2 \\sim\\exp \\left(- \\frac{3 M_p^4}{8V(\\phi)}\\right) . \\end{equation} Later this equation was also derived by Zeldovich and Starobinsky \\cite{ZelStar}, Rubakov \\cite{Rubakov}, and Vilenkin \\cite{Vilenkin} using the methods similar to the first method mentioned above (switching the sign of the Lagrangian). The corresponding wave function (\\ref{E35a}) was called ``the tunneling wave function.'' This wave function\\footnote{In fact, the two different ``derivations'' of this wave function described above lead to two slightly different wave functions \\cite{vilrecent}. However, since the difference between these two versions of the tunneling wave function is exponentially small, we will neglect it in this paper.} is dramatically different from the Hartle-Hawking wave function \\cite{HH}, as well as from the Vilenkin's wave function proposed few years earlier \\cite{NothVil}. An obvious objection against this result is that it may be incorrect to use different ways of rotating $t$ for the quantization of the scale factor and of the scalar field, see e.g. \\cite{HTnew}. If one makes the same rotation for for the matter fields as the rotation which we proposed for the scale factor, then one may encounter catastrophic particle production and other equally unpleasant consequences. On the other hand, as we have seen, if one assumes without any proof that it is enough to make the standard Wick rotation to quantize the scale factor, one does not obtain the wave function of the ground state $\\Psi_0$, and one gets the counterintuitive result that large universes are created much easier than the small ones. We believe that the problem here goes far beyond the issue of the Wick rotation. The idea that a consistent quantization of an unstable system of matter with positive energy density coupled to gravity with negative energy density can be accomplished by a proper choice of a complex contour of integration may be too optimistic. We know, for example, that despite many attempts to develop a Euclidean formulation of nonequilibrium quantum statistics or of the field theory in a nonstationary background, such a formulation still does not exist. It is quite clear from (\\ref{E32}) that the $t \\rightarrow -i\\tau$ trick does not give us the ground state wave function $\\Psi_0$ if the spectrum $E_n$ is not bounded from below. Absence of equilibrium, of any simple stationary ground state, seems to be a typical situation in quantum cosmology. A closely related instability is the basis of inflationary cosmology, where exponentially growing total energy of the scalar field appears as a result of pumping energy from the gravitational filed, whereas the total energy of matter plus gravitational field remains zero. Fortunately, in certain limiting cases this issue can be resolved in a relatively simple way. For example, at present the scale factor $a$ is very big and it changes very slowly, so one can consider it as a classical background, and quantize only the usual matter fields with positive energy. In this case one should use the standard Wick rotation $t \\rightarrow -i\\tau$\\@. On the other hand, in inflationary universe the evolution of the scalar field is very slow; during the typical time intervals $O(H^{-1})$ it behaves essentially as a classical field. Thus to a good approximation one can describe the process of creation of an inflationary universe filled with a homogeneous scalar field by the quantization of the scale factor $a$ only, and by the rotation $t \\rightarrow i\\tau$. When using the tunneling wave function, for example, for the description of particle creation in de Sitter space, instead of introducing a universal rule for the Wick rotation one should operate in a more delicate way, treating separately the scale factor and the particle excitations, see e.g. \\cite{VG}. Similarly, one should not use the Hartle-Hawking wave function for the description of creation of an inflationary universe, but one can use it for investigation of fluctuations of this background. These fluctuations are local, and often they appear simply as a result of quantum fluctuations of matter fields having positive energy. In particular, long-wavelength fluctuations of the scalar field $\\phi$ in inflationary universe may change local value of energy density $V(\\phi)$ inside the domains of a size greater than the size of the event horizon $H^{-1}$. For a comoving observer, such a change looks like a homogeneous change of the scalar field $\\phi$ and of the Hubble constant $H(\\phi)$, so he might want to (erroneously) interpret it as a result of quantum fluctuations of the scale factor. These are local perturbations of the homogeneous classical background. These perturbations are produced by fields with positive energy. Therefore in all situations where the inflationary background changes slowly (and in this sense can be considered a ground state of the system) one can use the Hartle-Hawking wave function for investigation of fluctuations of this background. For example, Hartle-Hawking wave function can be used for description of black hole formation in a pre-existing de Sitter background \\cite{BH}. But this method should not be used for description of quantum creation of de Sitter space with a pair of black holes in it. One can also obtain the amplitude of density perturbation in inflationary universe by a rather complicated method using the Hartle-Hawking wave function \\cite{HHAL}. However, the same results for density perturbations can be obtained by assuming that inflationary universe was created from nothing in accordance with the tunneling wave function, and then it expanded and produced perturbations in accordance to \\cite{Pert}. Moreover, as we already mentioned, in chaotic inflation there is no need to assume that any process of tunneling ever took place in the early universe. One may simply assume that the universe from the very beginning expanded classically, and then obtain the same results for the density perturbations using methods of Ref. \\cite{Pert}. Derivation of equations (\\ref{Vil1}), (\\ref{E36}) and their interpretation is far from being rigorous, and therefore even now it remains a subject of debate. From time to time this issue attracts a lot of attention. For example, the famous proposal to solve the cosmological constant problem in the context of the baby universe theory, which was very popular ten years ago, was based entirely on the use of the wrong sign of de Sitter action in the Hartle-Hawking approach to quantum gravity \\cite{b66,Coleman}. One of the main authors of this proposal, Sidney Coleman, emphasized: ``The euclidean formulation of gravity is not a subject with firm foundations and clear rules of procedure; indeed it is more like a trackless swamp. I think that I have threaded my way through it safely, but it is always possible that unknown to myself I am up to my neck in quicksand and sinking fast'' \\cite{Coleman}. After two years of intensive investigation of this issue it became clear that the wrong sign of the Euclidean action can hardly provide a reliable explanation for the vanishing of the cosmological constant. Moreover, recent observational data indicate that the cosmological constant may not vanish after all. To summarize, the derivation of the Hartle-Hawking wave function is rather ambiguous. Still, our main objection with respect to this wave function is related not to its derivation, but rather to its interpretation. The main purpose of the paper by Hartle and Hawking \\cite{HH} was to find the wave function describing the least exited, stationary state of the gravitational system, which would be analogous to the ground state on the harmonic oscillator or of the hydrogen atom. And indeed it gives a nice description of quantum fluctuations near de Sitter background, which in a certain sense is stationary. (There is a coordinate system where de Sitter space is static.) In such a situation one can consider matter fluctuations, and then find fluctuations of the scale factor induced by the fluctuations of matter. Then the problem of negative energy of the scale factor does not arise, and one can use the Hartle-Hawking wave function to study fluctuations in/of the pre-existing background. However, we do not see anything in the ``derivation'' of the Hartle-Hawking wave function which would indicate that it can be used for investigation of the probability of quantum creation of the universe. The tunneling wave function also has certain limitations, but it seems to have a better chance to describe the process of quantum creation of the universe. In the subsequent discussion an exact form of this wave function will not be important for us. The only property of this wave function which we are going to use is that quantum creation of the universe should not be strongly suppressed if it can be achieved by fluctuations of metric on the Planck scale $M_p^{-1}$ at the Planck density $M_p^4$. Since the debate concerning the wave function of the universe continues for the last 15 years, it may be useful to look at it from a somewhat different perspective, which does not involve discussion of ambiguities of the Euclidean quantum gravity. In the next section we will discuss the stochastic approach to quantum cosmology. Within this approach equations (\\ref{Vil1}) and (\\ref{E36}) can be derived in a much more clear and rigorous way, but they will have a somewhat different interpretation. ", + "conclusions": "Prior to the invention of the inflationary universe scenario it seemed that quantum cosmology is very important for understanding the underlying principles of the theory of evolution of the universe, but it may not have any observational consequences. During the last 15 years quantum cosmology has become a more established science, which allows us to make testable observational predictions. As we have seen, both the Hartle-Hawking and the tunneling wave function of the universe can describe creation of an open inflationary universe. This is a very interesting possibility in view of the recent tendency to claim that the observations favor smaller value of $\\Omega$. However, different versions of quantum cosmology predict completely different values of $\\Omega$. The Hartle-Hawking wave function predicts that if the universe is closed, then $\\Omega \\gg 1$, and if it is open, one has $\\Omega \\sim 10^{-2}$. This is experimentally unacceptable. In this paper we confirmed that this result is practically model-independent if galaxy formation occurs due to adiabatic density perturbations produced during inflation. One may try to avoid this conclusion by appealing to some unspecified versions of string theory or M-theory where the situation might be better \\cite{HT,HTnew}. But in the absence of any realization of this idea one may conclude that at the present time the Hartle-Hawking wave function, {\\it if used to calculate the probability of quantum creation of the universe}, is in a direct contradiction with observational data. Is it really possible to rule out the Hartle-Hawking wave function on the basis of these results? Perhaps such a conclusion would be premature. The main argument which pushed the most probable value of $\\Omega$ toward $10^{-2}$ was based on the equation for adiabatic density perturbations in a theory of a single scalar field, Eq. (\\ref{PROB2}). This conclusion can change if adiabatic perturbations are very small, and perturbations responsible for galaxy formation are isocurvature, or if they are produced by topological defects. For example, one may imagine that the phase transition which leads to the formation of topological defects occurs during the last stages of chaotic inflation, see e.g. \\cite{KL}. Then the defect production is a threshold effect, which occurs only if the universe is formed with a sufficiently large scalar field $\\phi$. In such a situation the Hartle-Hawking wave function will suggest that the scalar field should be as small as possible, but still large enough for the phase transition to take place, because density perturbations would be too small in the universe without strings. Then the unfortunate prediction $\\Omega = 10^{-2}$ may disappear, but it will be replaced by the fine-tuning of the moment of onset of the phase transition. Also, the possibility to produce the large scale structure of the universe using isothermal perturbations or topological defects is currently out of favor, so we are not sure whether one should consider it seriously. In our opinion, the whole problem appears here because one tries to apply the Hartle-Hawking wave function for the investigation of the probability of creation of the universe. Our analysis of this issue contained in Sections II and III suggests that it should not be used for that purpose. In particular, we have seen that stochastic approach to inflation unambiguously produces the same probability distribution as the Hartle-Hawking wave function, see Eq. (\\ref{E38a}). This equation has a simple interpretation: the Hartle-Hawking wave function (in agreement with its derivation in \\cite{HH}) describes the probability distribution to find the field $\\phi$ in a stationary state (if this state exists) {\\it after} the field relaxes towards the minimum of the effective potential. This wave function does not describe creation of the universe, inflation and the process of relaxation toward this ground state, which is the main subject of our investigation. If one uses the tunneling wave function for the description of initial conditions in the universe, then in most inflationary models the universe should have $\\Omega = 1$, which agrees with the standard expectation that inflation makes the universe flat. This result is not sensitive at all to the exact features of the tunneling wave function, and in fact to the very use of the tunneling wave function. The only thing which we need to assume is that there is no exponential suppression of quantum creation of a very small universe as compared to the probability of creation of a very large universe \\cite{book}. Moreover, according to the theory of a self-reproducing inflationary universe, which applies to most versions of chaotic inflation \\cite{b19}, one can avoid making even this assumption. The theory of a self-reproducing universe asserts that initial conditions are nearly irrelevant for the description of the properties of the main part of the universe \\cite{Mijic,LLM}. In most models of that type one has $\\Omega = 1$ after inflation. There exists a new potentially interesting class of models where creation of an open universe described by the tunneling wave function may be possible. A thorough investigation is needed in order to verify whether this possibility is realistic or not. There are many reasons to be sceptical about it, see Sect. \\ref{tunn} and also \\cite{VIL,BL}. It is important, however, that independently of this possibility we still have the class of models proposed in \\cite{Gott,BGT,Open}, which does not seem to work in the context of the Hartle-Hawking proposal, but which is quite compatible with the tunneling wave function of the universe, as well as with the theory of a self-reproducing inflationary universe. Investigation of quantum cosmology in application to the open universe creation is very difficult. Much work is to be done in order to investigate the new possibilities which we now have. However, one should not underestimate the recent progress. Until very recently, we did not have {\\it any} consistent cosmological models describing a homogeneous open universe. Even though the open universe model did exist from the point of view of mathematics, it simply did not appear to make any sense to assume that all parts of an infinite universe can be created simultaneously and have the same value of energy density everywhere. That is why it is very encouraging that during the last few years we have found several different mechanisms of creation of an open universe. All of these mechanisms require the universe to be inflationary. It is still true that inflationary models describing the universe with $\\Omega = 1$ are much simpler than the models with $\\Omega \\not = 1$. Hopefully, the universe will appear to be flat, and we will never need to use any of the models of open inflation. But if we find out that Nature has chosen to build the universe in a way which does not look particularly natural, this may give us a rare opportunity to reexamine some of our ideas and to learn more about quantum cosmology. \\subsection*" + }, + "9802/astro-ph9802280_arXiv.txt": { + "abstract": "We investigate the spectrum of photohadronically produced neutrinos at very high energies (VHE, ${\\gsim} 10^{14}\\eV$) in astrophysical sources whose physical properties are constrained by their variability, in particular jets in Active Galactic Nuclei (blazars) and Gamma-Ray Bursts (GRBs). We discuss in detail the various competing cooling processes for energetic protons, as well as the cooling of pions and muons in the hadronic cascade, which impose limits on both the efficiency of neutrino production and the maximum neutrino energy. If the proton acceleration process is of the Fermi type, we can derive a model independent upper limit on the neutrino energy from the observed properties of any cosmic transient, which depends only on the assumed total energy of the transient. For standard energetic constraints, we can rule out major contributions above $10^{19}\\eV$ from current models of both blazars and GRBs; and in most models much stronger limits apply in order to produce measurable neutrino fluxes. For GRBs, we show that the cooling of pions and muons in the hadronic cascade imposes the strongest limit on the neutrino energy, leading to cutoff energies of the electron and muon neutrino spectrum at the source differing by about one order of magnitude. We also discuss the relation of maximum cosmic ray energies to maximum neutrino energies and fluxes in GRBs, and find that the production of both the highest energy cosmic rays and observable neutrino fluxes at the same site can only be realized under extreme conditions; a test implication of this joint scenario would be the existence of strong fluxes of GRB correlated muon neutrinos up to ultra high energies, ${>}10^{17}\\eV$. Secondary particle cooling also leads to slightly revised estimates for the neutrino fluxes from (non-transient) AGN cores, which are commonly used in estimates for VHE detector event rates. Since our approach is quite general we conclude that the detection or non-detection of neutrinos above ${\\sim}10^{19}\\eV$ correlated with blazar flares or GRBs (\\eg, with the Pierre Auger Observatory), would provide strong evidence against or in favor of current models for cosmic ray acceleration and neutrino production in these sources. ", + "introduction": " ", + "conclusions": "" + }, + "9802/astro-ph9802249_arXiv.txt": { + "abstract": "Several classes of stars (most notably O and B main-sequence stars, as well as accreting white dwarfs and neutron stars) rotate quite rapidly, at spin frequencies greater than the typical g-mode frequencies. We discuss how rapid rotation modifies the $\\kappa$-mechanism excitation and observability of g-mode oscillations. We find that, by affecting the timescale match between the mode period and the thermal time at the driving zone, rapid rotation stabilizes some of the g-modes that are excited in a non-rotating star, and, conversely, excites g-modes that are damped in absence of rotation. The fluid velocities and temperature perturbations are strongly concentrated near the equator for most g-modes in rapidly rotating stars, which means that a favorable viewing angle may be required to observe the pulsations. Moreover, the stability of modes of the same $l$ but different $m$ is affected differently by rotation. We illustrate this by considering g-modes in Slowly Pulsating B-type stars as a function of the rotation rate. ", + "introduction": "\\label{sec:intro} Despite a long-standing interpretation of multiple periodicity in $\\beta$~Cepheids and Slowly Pulsating B-type (SPB, \\cite{Waelkens91}) stars (also referred to as 53 Per stars, after the prototype discovered by \\cite{smith-53per}) in terms of nonradial pulsations, the details of the excitation mechanism remained unresolved until quite recently. The driving of pulsations in the He$^+$ ionization zone, which gives rise to the classical Cepheid instability strip, cannot operate in these hot stars. In a remarkable example of the interplay between theory and observation, the updated opacity calculations (\\cite{OPAL92}; \\cite{OP94}) found a significant opacity increase in the metal partial ionization zone, just as suggested by \\cite{Simon82} as the explanation for intrinsic variability in 53 Per stars. The resulting instability region, determined by the $\\kappa$ mechanism operating in the ``metal opacity bump'' at $T\\approx 1.5\\times10^5\\ {\\rm K}$, incorporates most OB stars ($M\\gtrsim 3M_\\odot$) that are in the core hydrogen burning stage (\\cite{Dziembowski-bcep}; Dziembowski, Moskalik, \\& Pamyatnykh 1993; \\cite{GautschySaio-bstars}). The most recent excitation calculations for $\\beta$~Cepheids and SPBs (\\cite{Pamyatnykh-lanl}) place all known variable stars in these classes within the theoretical instability region. However, it is not clear that {\\it all\\/} stars in this region are indeed pulsating. All previous g-mode pulsation calculations have been carried out for non-rotating or very slowly rotating stars. However, most OB stars rotate relatively rapidly and observations indicate that g-mode pulsations and rotation are somehow related. As discussed in \\S\\ref{sec:freqs}, rotation strongly affects g-mode properties when $2\\Pmode\\gtrsim\\Pspin$, where $\\Pmode$ is the mode period in the {\\it co-rotating} frame, and $\\Pspin$ is the spin period. Rotation would have to be near breakup to affect the p-mode pulsations in the $\\beta$ Cepheids. However, typical SPB stars ($M\\approx 4 M_\\odot$, $R\\approx 3R_\\odot$) exhibit high radial order g-mode pulsations of day-like periods, so that rotational velocities \\begin{equation}\\label{eq:rot-condition} \\vrot\\ge 75 \\kms \\left(\\frac{1\\ {\\rm day}}{\\Pmode}\\right) \\left(\\frac{R}{3 R_\\odot}\\right), \\end{equation} will significantly affect the g-modes. The mean rotational velocity for B stars is $\\vrot\\approx200\\kms$ (\\cite{Lang-astro}) and thus cannot be neglected in pulsation calculations. One of the hallmarks of the photometrically discovered SPBs is that they are anomalously slow rotators, with projected rotational velocities $v \\sin i \\sim 5 - 40 \\kms$ (\\cite{Waelkens87}; \\cite{Waelkens91}). While most of the $\\sim70$ SPBs discovered by Hipparcos (\\cite{Waelkens98}) are also slow rotators, some have $v\\sin i$'s up to $150\\kms$, a regime where rotational effects are very important. Another correlation between rotation and pulsations is that no SPB stars were found in the open clusters NGC~3293 and NGC~4755, where the mean projected velocity of B stars is $v\\sin i\\approx170 \\kms$ (\\cite{Balona94b}). Nevertheless, there were many $\\beta$~Cep stars discovered {\\it photometrically\\/} in the same clusters (\\cite{Balona94a}; \\cite{BalonaKoen94}), so this was not due to a selection or metallicity effect. The fact that most photometrically discovered SPBs are slow rotators, as well as the absence of SPBs in these two clusters, led \\cite{Balona-review} to conjecture that rapid rotation suppresses g-modes in these stars. In this paper we explore the interaction between rapid rotation and nonradial pulsations. We begin in \\S\\ref{sec:freqs} by outlining the effects of rotation on the adiabatic mode properties (frequencies and eigenfunctions) and observability of g-modes in rotating B stars. In \\S\\ref{sec:excit}, we discuss the interplay between rapid rotation and excitation of pulsations by the $\\kappa$~mechanism and its effect on the stability of g-modes in B stars. We close by outlining the future directions for our work and mention other potential applications of the rotationally modified g-mode theory we are developing. ", + "conclusions": "\\label{sec:conclusions} Though we have made progress in understanding the interaction between rotation and pulsations, we have yet to unambiguously explain the preponderance of slow rotators among the field SPB stars. Rotational modification of the mode period can destroy the timescale match $\\Pmode\\sim\\tauion$, and stabilize a g-mode that is excited in a non-rotating SPB model. Conversely, if the thermal time at the driving zone $\\tauion$ does not match the non-rotating mode period, the mode can be ``spun up'' and become excited if the star is rotating rapidly. The periods of excited modes in the co-rotating frame are largely determined by the structure of the star, in order to satisfy $\\Pmode\\sim\\tauion$. However, the radial orders $n$ of the excited g-modes depend on rotation. In order to maintain the timescale match with increasing $\\Omega$, one needs to increase $n$ (note that the co-rotating frame periods and stability of g-modes with $m=-l$ is not significantly affected by rotation). However, going to a larger $n$ may mean more radiative damping in the stellar interior, especially in the $\\mu$-gradient zone. In addition, we find that different $m$ harmonics of the same mode are affected to various degrees by rotation, which may have consequences for line profile observations capable of determining the $l$ and $m$ values of the pulsation modes. Rotation does not significantly affect the g-modes in B stars until the spin period is comparable to the mode period. Rotation this rapid confines most g-modes to the equatorial region, requiring a favorable viewing geometry for detection, and can destroy the timescale match needed for mode excitation. In this sense, the correlation between slow rotation and SPB pulsations may be an observational selection effect after all, since the g-modes in rapid rotators, even if present, only cover a small fraction of the stellar surface. In order to differentiate between these two effects, one needs information about the inclination angle $i$, as well as the amplitude of pulsations. Though, as explained in \\S\\ref{sec:freqs}, rotation is never ``rapid'' (in the sense $2\\Omega\\gtrsim\\omega$) for p-mode pulsations, \\cite{LeeBaraffe95} have shown cases where the low-order rotational perturbations can affect p-mode stability in $\\beta$ Cep-like stars rotating close to breakup. Our work may also be applicable to other classes of variable stars, most notably the $\\gamma$~Doradus stars. These stars, named after the prototype discovered by \\cite{cou63}, are early F-type main-sequence stars that lie at or below the cool edge of the classical Cepheid instability strip, and show variability with periods of around a day (\\cite{bre95}). It is currently believed that the variability is due to g-modes of high radial order (\\cite{kri97}). For an F-type star with $R\\sim 1.5 R_\\odot$, rotational velocities that would affect g-modes with one-day period are $\\vrot\\ge 40\\kms$ (see eq.~[\\ref{eq:rot-condition}]). Typical rotational velocities for early F stars are $\\vrot\\approx100\\kms$ (\\cite{Lang-astro}). Therefore, if g-modes are indeed responsible for $\\gamma$~Dor variability, rotation will have a significant effect on the mode properties and stability." + }, + "9802/astro-ph9802139_arXiv.txt": { + "abstract": "Some recent work on radio source fields at $z\\sim 0.7$ and $z\\sim 4$ is discussed. At $z\\sim 0.7$ we find that radio-loud quasars are typically found in moderately rich environments independent of radio luminosity, consistent with previous results at $z\\sim 0.5$. In the field of the $z=4.41$ radio galaxy 6C0140+326 we find several candidate $z>4$ galaxies using the Lyman-break technique, two of which have detectable UV absorption features. In two $z>4$ radio galaxies, we find evidence for gravitational lensing affecting the fluxes by a few tens of percent, although we cannot rule out unusual lensing events which have larger magnifications associated with them. A simple calculation suggests lensing of $z>4$ radio sources should be very common. ", + "introduction": "The discoveries that powerful FRII radio galaxies and radio-loud quasars are frequently found in moderately rich clusters by $z\\sim 0.5$, in contrast to the situation at low redshifts, represent a large evolutionary change in the relatively recent past (Hill \\& Lilly 1991; Ellingson, Yee \\& Green 1991). Studying the cluster environments at higher redshifts is clearly important, and is possible in the optical up to $z\\sim 0.8$. At $z\\stackrel{>}{_{\\sim}}0.8$ when the 4000A break redshifts out of the optical window, cluster galaxies become hard to distinguish from the background. This can be partly improved by observing in the near-infrared, but even here the slow change of angular size distance with redshift means that clusters of a given physical size become no more compact on the sky. At the highest redshifts, spectral methods of identifying companion galaxies become necessary, for example narrow-band or, at sufficiently high redshift, Lyman-limit techniques. Foreground clusters may be important too. Claims of statistical correlations between luminous radio sources and foreground galaxies (Benitez et al.\\ 1997; Hammer \\& LeFevre 1990) are supported by correlations seen for different classes of AGN with foreground galaxy and cluster catalogues (e.g.\\ Bartsch, Barthelmann \\& Schneider 1997; Rodrigues-Williams \\& Hogan 1994) and by the detection of shear fields around a few $z\\sim 1$ radio sources (Fort et al.\\ 1996; Schneider et al.\\ 1997). The nature of the lensing population, and how it is able to produce these effects whilst retaining a plausible form for the AGN luminosity function remains a mystery, but large (i.e.\\ cluster-sized) lenses could lens both radio sources, any close companions, and any even higher redshift objects in the field. In this paper I begin by describing some early results of optical observations of the cluster environments of a sample of radio-loud quasars at $z\\sim 0.7$ made with the NOT, as part of an on-going programme to study the cluster environments of radio sources and radio-quiet quasars at $z\\sim 0.7$, in collaboration with Margrethe Wold and Per Lilje. I then discuss some work currently under way with my student Robin Stevens on $z>4$ Lyman limit systems, before discussing weak lensing results on the two most distant radio galaxies. ", + "conclusions": "" + }, + "9802/astro-ph9802005_arXiv.txt": { + "abstract": "s{The last decade has shown a considerable development of gravitational lensing for cosmology because it probes the amount and the nature of dark matter, and provides information on the density parameter $\\Omega$, the cosmological constant $\\Lambda$ and the Hubble constant $H_{o}$. Therefore, gravitational lensing can constrain the cosmological scenario which gave birth to the Universe as it appears today. The ongoing programs and future projects which are developing now all over the world show that gravitational lensing is considered as a major cosmological tool for the coming years as well. In this review, we summarize some of the most recent advances in the fields relevant for the dark matter issue. We will focus on the microlensing, the arc(let)s and the weak lensing studies. The possibility to check the existence of a non-zero $\\Lambda$ is presented elsewhere (see Fort et al. contribution). } ", + "introduction": "\\label{sec:intro} The present-day structuration of the Universe likely formed from gravitational condensations of primordial fluctuations. In a homogeneous and isotropic universe, the growth and the late evolution of these fluctuations depend on the amount of mass-energy presents in the Universe and the nature of its matter content. The former is described by the cosmological parameters $\\Omega$ and $\\Lambda$, whereas its nature can be inferred from the shape of the power spectrum of the initial fluctuations and the amount of baryonic matter we can observe today. These crucial quantities are then among the most challenging observing targets for the end of this century, and motivated also the launches of MAP and Planck-Surveyor by the beginning of 2000. \\\\ The large variety of observational techniques applied over a wide range of dynamical systems shows compelling evidence that most of them are dominated by invisible matter. Furthermore, it seems that the dark matter fraction increases with the mass range of the systems. The amount of dark matter deduced from these studies leads to the conclusions that \\ \\ (1) dark matter is the main component of the Universe, \\ \\ (2) the visible mass does not fully account for the baryonic mass permitted from the theoretical expectations of the Big Bang Nucleosynthesis (BBN), allowing part of the dark matter to be baryonic, \\ \\ (3) on the other hand, if the mass-to-light ratios inferred from observations are correct, dark matter cannot be only baryonic. \\\\ There is still room for controversy on these conclusions because the measurements of the amount and the distribution of the matter are {\\sl indirect}: \\ \\ (1) 2- and 3-dimension galaxy surveys only depict the distribution of {\\sl light}; \\ \\ (2) the mass of gravitational systems are not simply inferred because assumptions on the geometry of their mass and light distribution profiles and on their dynamical stage are necessary. Some of these hypotheses are much debated. \\ \\ (3) Finally, the dynamical studies of large-scale galaxy flows which map the large-scale mass distribution of the Universe are not yet conclusive because of the poorness of the catalogs. \\\\ In fact, the rotation curves of spiral galaxies seem to give by far the most robust estimates of the total mass of galaxies. Athough the mass-to-light ratios inferred for these galaxies requires that their halos are dominated by dark matter, the amount needed is compatible with the upper limit of the baryon fraction deduced from BBN, and does not require that halos have non-baryonic content. Hence, the search for the nature of dark matter in galaxies as well as for new robust mass estimators is important.\\\\ Gravitational lensing effects can directly probe deflecting masses and can determine without ambiguity the amount of matter present along the line-of-sight. Its astrophysical interest only raised after the discoveries of the first multiply imaged quasar \\cite{wcw}, the gravitational arcs \\cite{sfmp}$^{\\!,\\,}$ \\cite{smfmc}$^{\\!,\\,}$ \\cite{lp} and the arclets \\cite{fpmms}. But the on-going massive monitoring of microlensing events and the development of large programs for mapping the large-scale structures by using weak lensing make gravitational lensing effect one of the most promising tools to address some cosmological issues of the next ten years. This review only focus on the recent results relevant to dark matter issue. Section \\ref{sec:lensequa} summarizes the fundamental concepts and equations of gravitational lensing. In section \\ref{sec:baryongal}, the latest microlensing experiments are presented. Sections \\ref{sec:arclusters} and \\ref{sec:wlcluster} are devoted to investigations of clusters of galaxies and section \\ref{sec:wllss} to the promising investigation of large-scale structures. ", + "conclusions": "Microlensing surveys have shown that the existence of a baryonic halo dominated by brown dwarfs around our Galaxy is not confirmed. However, galaxy-galaxy lensing and Einstein rings suggest that galactic halos of dark matter are present. Instead, microlensing events reveal a new population of compact star-like objects may exist in our Galaxy. On the other hand, studies of arc(let)s and weak lensing in clusters show that $\\Omega>0.2$ is almost certain and provide evidence that $0.2<\\Omega<0.6$ on scales below 2.5 $h^{-1}$ Mpc. The next decade will provide important constraints on $\\Omega$ and $P(k)$, as well as first maps of mass-density of the Universe on 100$h^{-1}$ Mpc scales regardless the light distribution. This is an enthusiastic period for gravitational lensing applications to cosmology." + }, + "9802/astro-ph9802274_arXiv.txt": { + "abstract": "I present pointed ROSAT PSPC observations of the pre-cataclysmic binary V471~Tauri. The hard X-ray emission ($>$0.4\\,keV) is not eclipsed by the K star, demonstrating conclusively that this component cannot be emitted by the white dwarf. Instead I show that its spectrum and luminosity are consistent with coronal emission from the tidally spun-up K star. The star is more active than other K stars in the Hyades, but equally active as K stars in the Pleiades with the same rotation periods, demonstrating that rotation---and not age---is the key parameter in determining the level of stellar activity. Ths soft X-ray emission ($<$0.4\\,keV) is emitted predominately by the white dwarf and is modulated on its spin period. I find that the pulse-profile is stable on timescales of hours and years, supporting the idea that it is caused by opacity of accreted material. The profile itself shows that the magnetic field configuration of the white dwarf is dipolar and that the magnetic axis passes through the centre of the star. There is an absorption feature in the lightcurve of the white dwarf, which occurs at a time when our line-of-sight passes within a stellar radius of the K star. The column density and duration of this feature imply a volume and mass for the absorber which are similiar to those of coronal mass ejections of the Sun. Finally I suggest that the spin-orbit beat period detected in the optical by Clemens {\\it et al.} may be the result of the interaction of the K-star wind with the magnetic field of the white dwarf. ", + "introduction": "The eclipsing close binary V471~Tauri is a member of the Hyades open cluster, and contains a white dwarf and K2V star in a 12.5\\,h orbit \\cite{Nelson70,Young72}. The white dwarf is hot, $\\rm T=3\\times10^{4}\\,K$, and is a strong source of ultraviolet and soft-X-ray emission \\cite{Guinan84}. Observations with EXOSAT by \\scite{Jensen86} revealed a double-peaked modulation of the soft X-rays at a period of 555\\,s. They suggested this could be caused by either the changing viewing angle of bright and dark regions on the white dwarf, the 555\\,s period being the rotation period of the white dwarf, or radiatively-driven g-mode pulsations. Bright or dark regions could be formed by magnetic accretion of the K-star wind onto polar regions of the white dwarf: bright regions due to heating or dark regions due to opacity of accreted material. The detection of the same period in the optical \\cite{Robinson88,Clemens92}, but in anti-phase with the X-rays \\cite{Barstow92}, proved that the modulation could not be due to pulsations. It also showed that the polar regions of the white dwarf are probably dark in X-rays and bright in the optical. \\scite{Barstow92} presented ROSAT all-sky-survey observations of V471~Tau. They detected the white-dwarf spectrum but also some X-ray emission at energies higher than could be attributed to the white dwarf ($>$0.3\\,keV). Together with the observation that a fraction of the X-ray flux is not eclipsed \\cite{Jensen86}, this suggests that the K star is a significant source of X-rays. In this paper I present observations of V471~Tau made with the ROSAT position-sensitive proportional counter (PSPC) during the pointed phase of the mission. These confirm that the K star is a significant X-ray source, and show that its emission is consistent with that expected from a rapidly-rotating K star. The pointed ROSAT observations are the first with sufficient sensitivity and spectral resolution to separate the X-ray emission of the white dwarf and K star. ", + "conclusions": "\\subsection{The K star} I show conclusively, for the first time, that the hard X-ray component of V471~Tau is not eclipsed by the K star. I find that the spectrum and luminosity of this component is typical of rapidly-rotating K stars, and conclude that it is most-likely emitted by the K star. I note that this X-ray flux is normal for its rotation rate rather than its age (Fig.\\,\\ref{fig-xlum}). This demonstrates that rotation is the key parameter in determining stellar activity, and not age. In most studies, age and rotation rate are too closely related to be separable. I also observe an absorption dip in the X-ray emission of the white dwarf, absent in the K-star flux, which implies a column density and length-scale for the absorber similar to coronal mass ejection events seen on the Sun. I conclude that similar processes are probably at work on the K star in V471~Tau, and may account for the more dramatic orbital dips seen with EXOSAT and EUVE (because both these instruments are sensitive to a smaller absorbing column density). \\subsection{The white dwarf} \\label{sec-conc-wd} I detect the 555\\,s X-ray modulation discovered by \\scite{Jensen86}. The pulse-profile is identical to that measured with EXOSAT and EUVE, and is also stable throughout the ROSAT observations. This is consistent with the idea that the modulation is caused by opacity of material accreted at the magnetic poles of the white dwarf (since this opacity cannot vary on timescales shorter than the diffusion of metals in the atmosphere of the white dwarf). The profile itself is double-peaked, and the two peaks are separated in phase by precisely 180$^\\circ$. If the modulation is indeed caused by the opacity of accreted material, then the magnetic field of the white dwarf must direct the accretion flow onto two regions directly opposing each other. This shows that the magnetic field configuration must be dipolar, and that the magnetic axis must pass though the centre of the white dwarf. I do not detect the beat pulse discovered in the optical by \\scite{Clemens92} and place a 95\\% upper limit to its X-ray amplitude of 4\\%. Clemens {\\it et al.\\ } interpret this pulse as reprocessing of ionising radiation on the face of the K star. However the beat pulse is single peaked while the X-ray pulse is double peaked, and Clemens {\\it et al.\\ } can only reconcile these facts by assuming the X-ray emission does not trace the bulk of the ionising radiation. Instead I suggest that the beat pulse may reflect a modulation in accretion rate caused by the interaction of the wind of the K star with the magnetic field of the white dwarf. Beat periods arise naturally in systems where the accretion flow has memory of orbital phase \\egcite{Warner86}. However the same problem arises with this interpretation because in a high-inclination system, such as V471~Tau, where we see both magnetic poles, one would expect to see modulation at the frequency $\\rm (2\\,f_{spin}-f_{orbit})$ rather than that detected by Clemens {\\it et al.\\ } $\\rm (f_{spin}-f_{orbit})$ \\cite{Wynn92}. Still, as the geometry of magnetic wind accretion is poorly known, there may be scope to explain the single peaked beat pulse." + }, + "9802/gr-qc9802005_arXiv.txt": { + "abstract": "We consider compensated spherical lens models and the caustic surfaces they create in the past light cone. Examination of cusp and crossover angles associated with particular source and lens redshifts gives explicit lensing models that confirm the claims of Paper I \\cite{bi:I}, namely that area distances can differ by substantial factors from angular diameter distances even when averaged over large angular scales. `Shrinking' in apparent sizes occurs, typically by a factor of 3 for a single spherical lens, on the scale of the cusp caused by the lens; summing over many lenses will still leave a residual effect. ", + "introduction": "\\label{sec: Intro} This paper continues the study started in Paper I \\cite{bi:I} of how area distances behave in universes where strong gravitational lensing takes place. That paper considered the claim \\cite{bi:weinberg} that although individual lensing masses alter area distances for ray bundles that pass near by them, photon conservation guarantees the same area distance-redshift relation as in exact It was shown in \\cite{bi:I} that this claimed compensation result is incorrect once one has passed caustics, which are necessarily the result of strong gravitational lensing; consequently (by continuity) the result is not true in general. Indeed it has to be wrong because area distances are determined by the gravitational field equations (essentially through the null Raychaudhuri equation) quite independently of the issue of photon conservation (which is determined by Maxwell's equations and is valid whatever the space-time curvature). Thus the latter cannot causally determine area distances. In fact at large distances, `shrinking' takes place in that distant areas subtend smaller solid angles than they would in a FL universe model; and this effect will remain even when the observations relate to large angles. The way these small effects for individual lenses add up to give a significant averaged effect over the whole sky is discussed in Paper I. This may affect high-redshift number counts and Cosmic Background Radiation (`CBR') anisotropy observations at very small angles. The general argument has been given in I. Specific spherically symmetric examples (viewed from the centre, and so without caustics) are presented in \\cite{bi:II}, which thereby gives a rigorous proof of the existence of the effect we claim; but the models used are unrealistic as models of the real universe. In this Paper (Paper II), we consider exact examples with caustics displaying the shrinking effects discussed in Paper I. We show how to calculate the magnitude of the effect analytically and numerically for single spherically symmetric compensated lenses to which we can apply the thin-lens approximation, and look in detail at `top hat' lenses, which are the simplest in this class. We refer to and mainly follow the notation of Schneider, Ehlers and Falco (1992) \\cite{bi:Schneider92} (`SEF'). ", + "conclusions": "\\label{sec: concl} Of particular interest is the way the cusp size and the ``shrinking\" vary with redshift of the source and of the lensing object. This depends on two things: firstly the variation of angular sizes with redshift, remembering (a) minimum angular apparent diameter occurs at $z = 1.25$, so that the maximum angle $\\theta_c$ for cusps to form due to lenses of given size and strength will have minimum at that redshift; and (b) that the ratio of distances that enters $ \\sigma_{cr}$ saturates with increasing $z_s$ (for given $z_d$) but has a maximum for each $z_s$ at a $z_d$ of about $0.6$ which is thus the optimal distance for the lens in order to create cusps on the last scattering surface. that are not typical of all galaxies or clusters; but they confirm in a concrete way the broad picture proposed in Paper I: an area `shrinking' factor of 3 will occur for each lens that causes cusps, on the scale of the cusps (precisely: at the cut-off angle $\\theta_3$ which gives the same deflection at the source as the cusp angle, but on the opposite side). The total effect when averaging over large angular scales will depend on what fraction of the sky is covered by these angles for all lenses at all smaller scales, as a function of redshift; some simple estimates of this overall effect were given in Paper I. To determine realistic multiplicity factors as a function of redshift will require simulations with multiple lensing and more realistic lens models, for example standard elliptical lens models determined by a velocity dispersion parameter and ellipticity parameters as in \\cite{bi:Kneib} which allow an increase in the degree of multiple covering (because individual elliptic lenses can have a covering factor of 5). The effect will differ on angular scales, and will almost certainly be substantial due to micro-lensing, with an additional increase due to galactic and cluster lensing. The implication of this paper and Paper I is that it is incorrect to assume that areas average out to the background FL value on large angular scales; one can only know the true area ratios - expressed in the shrinking factors considered in this paper - by detailed calculation." + }, + "9802/astro-ph9802042_arXiv.txt": { + "abstract": "The SW~Sextantis stars are a group of cataclysmic variables with distinctive observational characteristics, including absorption features in the emission line cores at phases 0.2--0.6. Hellier and Robinson have proposed that these features are caused by the accretion stream flowing over the accretion disk. However, in a simple model the absorption occurred at all orbital phases, which is contradicted by the data. I show that invoking a flared accretion disk resolves this problem. ", + "introduction": "The SW~Sextantis stars are a subclass of the novalike variables, which are themselves a subclass of the cataclysmic variables (CVs). The qualification for novalike status is a stable accretion disk, presumably due to a mass transfer rate sufficiently high to prevent the instabilities associated with dwarf novae (e.g.~Warner 1995). The SW Sex stars are novalikes showing all or most of the following properties (e.g.~Thorstensen \\etal\\ 1991): (i) single-peaked emission lines, particularly \\HeIIl, incompatible with an origin in a Keplerian disk; (ii) gross asymmetries in the emission lines from the disk, so that they do not reflect the orbital motion of the white dwarf, (iii) peculiar `phase 0.5' absorption features in the core of the line during the orbital phases \\sqig0.2--0.6; (iv) a tendency to have orbital periods in the range 3--4 hr, just above the CV period gap; and (v) a high probability of being eclipsing (this one is most likely a selection effect). Many ideas have been proposed to explain SW~Sex stars, the favourites being strong accretion disk winds (Honeycutt \\etal\\ 1986; Dhillon \\etal\\ 1991; Hellier 1996); magnetically controlled accretion (Williams 1989; Casares \\etal\\ 1996); and mass-transfer streams which penetrate or flow over the disk (Shafter \\etal\\ 1988; Hellier and Robinson 1994). This paper addresses the combination of those ideas proposed in Hellier and Robinson (1994; hereafter Paper~1) and Hellier (1996; Paper~2), although for alternative views and discussion of other models see Dhillon \\etal\\ (1991), Casares \\etal\\ (1997) and Hoard and Szkody (1997). We suggested that SW~Sex stars are novalikes with abnormally high mass-transfer rates. This causes, first, strong winds from the inner disk or boundary layer, explaining the single-peaked line profiles. Second, it allows the accretion stream to overflow the initial impact with the edge of the disk rim, and continue on a free-fall trajectory to a second impact much nearer the white dwarf. Line emission from this re-impact causes the highly asymmetric line profiles. Further, we proposed that the stream is seen in absorption between the initial impact with the disk and its re-impact. This means that the normal emission `S-wave' common in other CVs is absent, or even in absorption. Also, we noted that the absorption from the stream could explain the phase 0.5 absorption features of SW~Sex stars. The remaining characteristic, the concentration on orbital periods just above the period gap, presumably results from evolutionary effects driving systems with those periods to higher mass transfer rates (\\mdot; e.g.~Shafter 1992). Indeed, in Paper~2 we noted that SW~Sex stars often show VY~Scl low states, implying that the abnormally high \\mdot\\ in the high state is balanced by periods of much lower \\mdot. This could result from irradiation-driven feedback cycles in \\mdot\\ (e.g.~Wu \\etal\\ 1995). In Papers 1 and 2 we computed the velocities expected from an accretion stream flowing over the disk, as a function of orbital phase, and turned these into model line profiles for comparison with data from PX~And and V1315~Aql. Overall the simulations supported the model, showing that the overflowing stream had the right velocity variations to explain the distorted line profiles and the phase 0.5 absorption. However, these simulations had a fundamental limitation in that they calculated velocities only, making no allowance for obscuration of one component by another, and so could not reproduce variations in the strength of features round the orbital cycle. Thus the simulations contained `phase 0.5 absorption' at all orbital phases, whereas the data doesn't. Since several authors (e.g.~Casares \\etal\\ 1996; Hoard and Szkody 1997) have cited this flaw as a primary reason for doubting the model, I have re-written the simulation code to include all spatial and obscuration effects (it is still a simple geometric model, though, making no attempt at radiative transfer). I pick up a suggestion from Paper~1 that disk flaring can confine the absorption to a limited phase range. This gives a much closer resemblance to the data, solving the biggest discrepancy between the data and models from Papers 1 and 2. ", + "conclusions": "The `phase 0.5 absorption' features seen in SW~Sex stars can be explained by a proportion of the accretion stream flowing over the disk. The fact that absorption is seen only at phases 0.2--0.6, and not at every phase as predicted by simulations in Papers 1 and 2, is explained if the accretion disk is flared at an angle of $\\approx$ 4\\deg. The simulation code has been re-written to include all spatial and obscuration effects, and the results in Papers 1 and 2 are still valid. If disk-overflow is a generally correct model for SW~Sex stars the degree of overflow could still vary with time or between systems. A reduction in disk-overflow would leave an SW~Sex star looking much like a normal novalike (as seems to have been the case for SW~Sex when observed by Dhillon, Marsh \\&\\ Jones 1997). In a high inclination novalike the splash caused by the stream--disk impact can itself produce absorption when it is at inferior conjunction, which would produce absorption dips centered on phase \\sqig 0.8, rather than the phase 0.2--0.6 characteristic of the SW~Sex stars. Dips at phase \\sqig 0.8 are seen in novalikes such as TV~Col (Hellier 1993) and BP Lyn (Hoard \\&\\ Szkody 1997)." + }, + "9802/astro-ph9802104_arXiv.txt": { + "abstract": "We present the light curves obtained during an eight-year program of optical spectroscopic monitoring of nine Seyfert 1 galaxies: 3C 120, Akn 120, Mrk 79, Mrk 110, Mrk 335, Mrk 509, Mrk 590, Mrk 704, and Mrk 817. All objects show significant variability in both the continuum and emission-line fluxes. We use cross-correlation analysis to derive the sizes of the broad \\Hbeta-emitting regions based on emission-line time delays, or lags. We successfully measure time delays for eight of the nine sources, and find values ranging from about two weeks to a little over two months. Combining the measured lags and widths of the variable parts of the emission lines allows us to make virial mass estimates for the active nucleus in each galaxy. The virial masses are in the range $10^{7-8}$\\,\\Msun. ", + "introduction": "It has been known since the late 1960s that both the continuum (e.g., Fitch, Pacholczyk, \\& Weymann 1967) and broad emission lines (e.g., Andrillat \\& Souffrin 1968) in active galactic nuclei (AGNs) vary in flux with time. In the 1980s, spectroscopic monitoring programs showed that the continuum and emission-line variations are closely coupled, confirming that the emission-line regions are powered predominantly by photoionization by the central source. However, the typical emission-line response times were found to be surprisingly short compared to the light-travel times expected by most photoionization equilibrium models (e.g., Peterson et al.\\ 1985). Attempts were made to determine light-travel times for the broad-line region (BLR) by cross-correlation of continuum and emission-line fluxes (e.g., Gaskell \\& Sparke 1986), but the analyses were plagued by sparsely sampled light curves and relatively large uncertainties in the measured fluxes (e.g., Gaskell \\& Peterson 1987; Edelson \\& Krolik 1988). In spite of the difficulties, the potential for extracting physical information about the central regions of AGNs from variability was generally regarded to be enormous (see Peterson 1988 for a review of the early monitoring programs and their implications): in principle, it is possible to constrain significantly the structure and kinematics of the BLR by determining the emission-line response to continuum variations as a function of wavelength, since the broad-lines are well-resolved in radial velocity even at low ($\\sim10$\\,\\AA) resolution. This process is known as ``reverberation mapping'' (Blandford \\& McKee 1982). Late in the 1980s, it became possible to obtain the quality and quantity of data necessary to determine emission-line response times. In the ultraviolet, large amounts of {\\it International Ultraviolet Explorer}\\ time were devoted to AGN variability projects (e.g., Clavel et al. 1991). In the optical, CCDs became widely available on even moderate-size ($\\sim2$-m) and small-size ($\\sim1$-m) ground-based telescope, making it possible to obtain high signal-to-noise ratio ($S/N$) spectra of high photometric accuracy with relative ease. The problem of poor time sampling was obviated by cooperation between observers, either by using many telescopes (e.g., Peterson et al.\\ 1991) or by a group using a single facility (e.g., Maoz et al.\\ 1990, Robinson 1994). Progress in reverberation mapping through 1992 is reviewed by Peterson (1993), and a more recent review is given by Netzer \\& Peterson (1997). In 1988, we began a long series of approximately weekly spectroscopic monitoring of nearby bright Seyfert galaxies with a CCD spectrograph on the 1.8-m Perkins Telescope at Lowell Observatory. We present here the first analysis of most of these data. The scientific goals of the program have been: \\begin{enumerate} \\item To acquire optical continuum and emission-line light curves of sufficient sampling and quality to determine accurately the emission-line response times, or ``lags'' for a number of AGNs. \\item To investigate AGN continuum behavior over a long temporal baseline. \\item To investigate the nature of broad emission-line {\\em profile}\\ variability, and see what this reveals about the kinematics of the broad-line region. \\item To investigate the possibility of structural changes in the BLR on time scales of years (which corresponds to the dynamical time for the BLR). \\end{enumerate} Some of the results of this program have been reported elsewhere in the literature, but this is the first comprehensive presentation of the data obtained since 1988. In this paper, we present the light curves for the optical continuum and the broad \\Hbeta\\ emission lines in nine Seyfert galaxies. For most of these objects, we are able to determine accurately the emission-line lags, and for these objects we can estimate the mass of the central black hole. In future papers, we will discuss other issues, such as line-profile variations. In \\S\\,2 we describe the observations and data reduction that led to this homogeneous data base of spectra. Analysis of the light curves is described in \\S\\,3, and we summarize our results in \\S,4. ", + "conclusions": "In this contribution, we have reported on the initial results of an eight-year spectroscopic monitoring program on Seyfert 1 galaxies. Most of these sources were observed nearly weekly whenever they were accessible, making these and NGC 5548 (Peterson et al.\\ 1994) especially well-suited to the study of long-term continuum and emission variability in non-blazar AGNs. For the best-sampled galaxies in our sample, we have over 100 homogeneous spectra. Of the nine sources presented in this paper, we were able to measure \\Hbeta\\ response times, or lags, for eight of them. The lags range from a little more than two weeks to more than two months, as summarized in Table 7. In many cases, the lags are measurable with data from individual observing seasons. These sometimes show variations from year to year, and this probably indicates that the BLR is physically thick, i.e., the outer radius of the BLR is much greater than the inner radius. We have combined the measured BLR response times with measured widths of the \\Hbeta\\ rms profiles to obtain virial mass estimates for the central source. These are all in the range $10^{7-8}$\\,\\Msun. We are grateful for support of this program by the National Science Foundation under grants AST--9420080, and its predecessors, AST--8702691, AST--8915258, and AST--9117086. This extensive program was made possible through the kind cooperation of our colleagues at Ohio State University and Lowell Observatory. We thank the current and past Directors of Lowell Observatory, R.L.\\ Millis and J.S.\\ Gallagher, for their support of this project for so many years. Maintaining the CCDS for weekly observations for nearly a decade has been successful on account of the capable work of B.\\ Atwood, P.L.\\ Byard, K.\\ Duemmel, A.A.\\ Henden, J.A.\\ Mason, T.P.\\ O'Brien, and R.J.\\ Truax at Ohio State and R.\\ Nye and R.\\ Oliver at Lowell Observatory, and we gratefully acknowledge their contributions to this program. We thank several Ohio State students for their participation in this project: B.\\ Ali, T.M.\\ Kassebaum, K.T.\\ Korista, N.J.\\ Lame, P.A.\\ Popowski, S.M.\\ Smith, and R.J.\\ White. We also thank M.\\ Dietrich for helpful comments and corrections. \\clearpage" + }, + "9802/astro-ph9802348_arXiv.txt": { + "abstract": "\\noindent One of the most dramatic discoveries made so far with the {\\em Rossi X-Ray Timing Explorer\\/} is that many accreting neutron stars with weak magnetic fields generate strong, remarkably coherent, high-frequency X-ray brightness oscillations. The \\hbox{$\\sim$325--1200~Hz} quasi-periodic oscillations (QPOs) observed in the accretion-powered emission are almost certainly produced by gas orbiting very close to the stellar surface and have frequencies related to the orbital frequencies of the gas. The \\hbox{$\\sim$360--600~Hz} brightness oscillations seen during thermonuclear X-ray bursts are produced by one or two hotter regions on the stellar surface and have frequencies equal to the stellar spin frequency or its first overtone. Measurements of these oscillations are providing tight upper bounds on the masses and radii of neutron stars, and important new constraints on the equation of state of neutron star matter. ", + "introduction": "Since the birth of X-ray astronomy 35 years ago, scientists have sought to use the X-radiation that comes from near the event horizons of black holes and the surfaces of neutron stars to probe quantitatively the strong gravitational fields near these objects and to determine the fundamental properties of dense matter (see, e.g., \\cite{ELP86,LP79}). The {\\em Rossi X-Ray Timing Explorer\\/} (\\rxte), which was launched on December 30, 1995, was specially designed to have the large area, microsecond time resolution, high telemetry bandwidth, and pointing flexibility needed to address these questions (see \\cite{Swank95}). With \\rxte, strong, high-frequency X-ray brightness oscillations have been discovered from at least two black holes (see \\cite{McClintock98}) and sixteen neutron stars (see \\cite{MLP98a,vdK98}). As a result, we appear to be on the threshold of achieving this decades-old goal of X-ray astronomy. Here we focus on the oscillations discovered in neutron stars and their implications for the properties of these stars and for neutron star matter. ", + "conclusions": "The discovery using the {\\em Rossi X-Ray Timing Explorer\\/} that many neutron stars with weak magnetic fields produce strong $\\sim$300--1200~Hz X-ray brightness oscillations is a spectacular achievement that validates both the scientific expectations that led to the mission and the long years of hard work that were needed to bring it to fruition. The kilohertz QPOs discovered in the accretion-powered emission are already providing interesting new upper bounds on the masses and radii of neutron stars, and on the stiffness of neutron star matter. The high-frequency oscillations discovered in the emission during thermonuclear X-ray bursts are likely to provide interesting new bounds on the compactness of neutron stars and hence on the softness of neutron star matter. Observation of a QPO with a frequency just 100~Hz higher than the highest frequency so far seen would exclude the stiffest proposed neutron star matter equations of state. Observation of innermost stable circular orbits would be the first confirmation of a strong-field prediction of general relativity and would fix the mass of the star involved, for each equation of state considered. Although there is currently no strong evidence that an innermost stable circular orbit has been discovered around any of these neutron stars, there is reason to hope that such evidence may be forthcoming. Given the rapid pace of discoveries with \\rxte, the prospects for obtaining compelling evidence of an innermost stable circular orbit appear good. \\medskip This work was supported in part by NSF grant AST~96-18524, NASA grant NAG~5-2925, and NASA RXTE grants at the University of Illinois, and by NASA grant NAG~5-2868 at the University of Chicago." + }, + "9802/astro-ph9802038_arXiv.txt": { + "abstract": "% We model the thermal X-ray profiles of Geminga, Vela and PSR 0656+14, which have also been detected as $\\gamma $-ray pulsars, to constrain the phase space of obliquity and observer angles required to reproduce the observed X-ray pulsed fractions and pulse widths. These geometrical constraints derived from the X-ray light curves are explored for various assumptions about surface temperature distribution and flux anisotropy caused by the magnetized atmosphere. We include curved spacetime effects on photon trajectories and magnetic field. The observed $\\gamma $-ray pulse profiles are double peaked with phase separations of 0.4 - 0.5 between the peaks. Assuming that the $\\gamma $-ray profiles are due to emission in a hollow cone centered on the magnetic pole, we derive the constraints on the phase space of obliquity and observer angles, for different $\\gamma $-ray beam sizes, required to produce the observed $\\gamma $-ray peak phase separations. We compare the constraints from the X-ray emission to those derived from the observed $\\gamma $-ray pulse profiles, and find that the overlapping phase space requires both obliquity and observer angles to be smaller than $20 - 30^0$, implying $\\gamma $-ray beam opening angles of at most $30-35^0$. ", + "introduction": "}% The multifrequency observations of $\\gamma $-ray pulsars may potentially constrain the geometry and location of emission regions in a neutron star (NS) magnetosphere. Such a constraint is necessary for our understanding of the entire picture of pulsar emission in different energy bands. However, in practice, the interpretation of high-energy observations is rather ambiguous and involves a number of model assumptions. The recent X-ray observations of pulsars Geminga, Vela, and PSR 0656+14 seem to indicate that the pulsed emission has a two-component X-ray spectrum (\\cite{oge95}, \\cite{hal97}, \\cite{str97}), where the soft (pulsed) component is likely due to thermal emission. Theoretical work by Pavlov et al. (1994) has shown that X-ray pulsation reaching the observed pulsed fractions can result from anisotropic emission in a magnetized atmosphere, even when the entire NS surface radiates. If this interpretation is correct, then the peak in the soft thermal X-ray profile is near the phase of closest approach to the magnetic pole. It is thus interesting to explore the relation between X-ray profiles as thermal black-body emission from the whole stellar surface ``beamed'' along the magnetic field in a NS atmosphere and non-thermal pulsed $\\gamma $-ray emission from these pulsars generated in the innermost magnetosphere of a NS (polar cap model)(Harding \\& Muslimov 1997a). The polar cap models for pulsed gamma-ray emission (Daugherty \\& Harding 1996, Sturner \\& Dermer 1994) predict double-peaked pulse profiles, where the closest approach to the pole is centered between the pulses. The peak of the broad thermal X-rays profile should thus also occur between the two gamma-ray pulses. In this paper we model the soft X-ray and $\\gamma$-ray light curves for the pulsars Geminga, Vela, and PSR 0656+14 surveying all possible orientations (see Figure 1) between the magnetic and spin axes (referred to as the {\\bf obliquity} angles) and the angle between the observer's line of sight and spin axis (referred to as the {\\bf observer} angles). In our modeling of the X-ray and $\\gamma $-ray light curves we assume that the NS has a centered dipole-like surface magnetic field, and that the thermal flux from the stellar surface corresponds to that expected in a cooling NS of age $\\sim 10^4-3\\times 10^5$ yr (the age category of Geminga, Vela and PSR 0656+14). In our analysis of the main observational features we consciously avoid additional model assumptions such as the possibility of polar cap heating by the precipitating relativistic particles and $\\gamma$-rays, off-setting of the magnetic axis of the NS and/or the presence of higher-order multipoles on the stellar surface (which are not unreasonable at all for future modeling), and occurrence of any specific cooling scenario (e.g. such as those with internal heating of the NS, or those allowing for any of the countless variants of rapid or slow cooling, etc). Also, we have not considered very compact NS models for which the effects of strong gravity by themselves may result in quite interesting signatures (\\cite{shib95}). All these possibilities, being attractive for a theoretical study, are rather abstract when discussed in the context of the currently available X-ray and $\\gamma $-ray observational data on pulsars. The quality of the observations and also the complexity of theoretical modeling (which usually involves many free parameters and poorly justified model assumptions) do not allow any conclusive statement regarding any of the abovementioned possibilities. The main goal of our study is to demonstrate that a polar cap model for the interpretation of the $\\gamma $-ray emission is viable at least for Geminga, Vela, and PSR 0656+14, and that the main observed X-ray and $\\gamma $-ray pulse characteristics (X-ray pulsed fraction, half-width of the X-ray profile, phase separation of $\\gamma $-ray peaks, and a phase shift $\\sim 0.1 - 0.2$ between the $\\gamma $-ray peaks) for these pulsars can be understood within the framework of a standard NS model with a dipole-like magnetic field and a relatively small obliquity ($\\lsim 30-35^o$). We begin our paper (\\S~2) with a summary of the observational X- and $\\gamma $-ray data on pulsars Geminga, Vela, and PSR 0656+14. In $\\S ~3$ we describe the details of our modeling of the observed soft X- and $\\gamma $-ray emission, and in $\\S ~4$ we present the results of our numerical calculations. Our principal conclusions are summarized in $\\S ~5$. ", + "conclusions": "}% In Figure 2 we show soft X-ray and high-energy $\\gamma$-ray pulse profiles for Geminga, Vela and PSR 0656+14. We have chosen these sources for modeling X-ray and $\\gamma$-ray pulsed profiles because they have been identified as having well-defined thermal components. Recent hard X-ray observations of Geminga (Halpern \\& Wang 1997) and PSR 0656+14 (Greiveldinger et al. 1996) with ASCA and Vela with RXTE (Rossi X-ray Timing Explorer; Strickman, Harding \\& De Jager 1998) have clearly defined the existence of separate non-thermal components and therefore greatly strengthened the interpretation of the soft X-ray components seen by ROSAT as thermal in origin. ROSAT observations of these sources had revealed phase shifts between the pulses seen in low energy ($0.1 - 0.5$ keV) and high energy bands ($0.5 - 2.0$ keV). Better definition of the pulse profile and spectrum in the energy range 1.0 - 30. keV by ASCA and RXTE have shown that, in the case of Vela and Geminga, the high energy pulses are double-peaked and in phase with the $\\gamma$-ray pulses measured by EGRET. Furthermore, the 2.0 - 30 keV spectrum of Vela is consistent with an extrapolation of the OSSE spectrum (Strickman et al. 1998). The characteristics of the hard X-ray components in these pulsars are therefore best explained if their origin is non-thermal magnetospheric emission. The phase shifts between the hard and soft components seen by ROSAT may be understood as a transition from a single broad thermal profile to a double-peaked non-thermal profile. The measured pulsed fractions and pulse widths are strongly energy dependent. The pulsed fraction (defined in equation~[\\ref{fp}]) of Geminga and PSR 0656+14 increases through the ROSAT band, starting at about $10-20\\%$ around 0.1 keV and reaching $80-90\\%$ above 1 keV. But it is not clear how much of this increase is instrinsic to the thermal component and how much is due to contamination by the hard, non-thermal component, which is known to have a high pulsed fraction. In the case of Geminga (Halpern \\& Wang 1997), the pulse profile also changes significantly above about 0.5 keV, where the power law spectral component becomes significant, indicating pulsed fraction contamination by the non-thermal component at the higher ROSAT energies. We have therefore chosen to model the thermal X-ray pulse profiles only in the lowest energy ROSAT bands available. For the purposes of our modeling, we have taken the observed pulse fractions of $20-30\\%$ for Geminga, based on the observed profiles given by Halpern \\& Wang (1997) for the range 0.08 - 0.28 keV. Table 1 summarizes the observed parameters of the X-ray and $\\gamma$-ray emission that are relevant to our modeling. The observed pulsed fractions for Vela and PSR 0656+14 are 11\\% (\\cite{oge93}) and 7\\% (\\cite{fin92}, \\cite{fin97}, \\cite{and93}), respectively. The pulse widths (FWHM -- Full Width at Half Maximum) in the lowest energy bands for all these pulsars lie roughly between 0.35 and 0.5 of pulse phase. We have determined approximate pulse widths, from the definition given in $\\S$ 4, of $0.35 \\pm 0.05$ for Geminga, $0.45 \\pm 0.05$ for Vela, and $0.55 \\pm 0.05$ for PSR0656+14. Vela and Geminga have well-defined high-energy $\\gamma$-ray profiles, as seen by EGRET, showing two sharp peaks with phase separations of 0.4 and 0.5 respectively and emission between the peaks (\\cite{Kan94}, \\cite{MH94}). PSR 0656+14, however, is considered only as a possible detection of pulsed emission by EGRET, with a poorly defined $\\gamma$-ray profile (\\cite{RM96}). For this study, we will assume a double-peaked profile with phase separation of 0.4, but we emphasize that this is highly uncertain. The relative phases of the X-ray and $\\gamma$-ray profiles are reasonably well determined for Vela and Geminga, but very uncertain for PSR 0656+14. The absolute phase was not determined at the time of the ROSAT measurement (\\cite{fin92}), but we have taken in Figure 2 the relative phase between the EGRET and soft ROSAT profiles as determined by Thompson (1997) from comparison of EGRET, ASCA (both having absolute timing) and hard ROSAT profiles. For all three pulsars, the broad X-ray profiles roughly coincide in phase with the $\\gamma$-ray profile. The peaks in the X-ray profile lie between the $\\gamma$-ray peaks in the case of Geminga and Vela, although both are offset by about 0.1 in phase toward the second $\\gamma$-ray peak. }% We have discussed the issue of whether the observed $\\gamma $-ray emission and recently detected soft (presumably thermal) X-ray emission from Geminga, Vela and PSR 0656+14 can be understood in terms of a polar cap (or hollow cone) model proposed for the $\\gamma $-ray emission from pulsars. We have calculated the range of observer and obliquity angles allowed by the observed pulsed fractions and widths of the soft X-ray profiles of these pulsars (at the median energy of $\\sim $0.18 keV), assuming anisotropy of the X-ray emission in a magnetized atmosphere. We found that small values of observer angle and obliquity are required to account for the single, relatively broad (with a phase width of $\\sim $0.35-0.5) X-ray peaks and these values can still produce the observed pulsed fractions of the X-ray profiles. The range of these angles restricted by the X-ray profiles are found to be consistent with those values required to reproduce the observed narrow double $\\gamma $-ray peaks separated by a phase interval of 0.4-0.5. In addition, the appearance of a single broad X-ray pulse between the two $\\gamma$-ray peaks predicted by polar cap models seems to be borne out at least for Vela and Geminga. Our main results can be summarized as follows. \\begin{enumerate} \\item The possibility of beaming of the thermal X-ray emission in Geminga, Vela, and PSR 0656+14 provides a consistent explanation for their observed X-ray light curves and is in accord with the polar cap models for their $\\gamma $-ray emission. \\item The obliquity and/or observer angle in Geminga, Vela, and PSR 0656+14 may be less than 30$^0$. The $\\gamma$-ray opening angles must be at least 13$^0$ for Geminga and at least 5$^0$ for Vela and PSR 0656+14. \\item For anisotropic X-ray emission, the observed pulsed fraction and pulse width are much less sensitive to the effective temperature and are determined primarily by the degree of a beaming. \\item The magnitude of X-ray pulsed fraction is mainly determined by the magnitude of the beamed emission relative to the fan beam emission. In the case of a very strong beamed component, whose contribution to the X-ray flux is significant, the maximum pulsed fractions should be observed for a rather wide range of obliquity and observer angles. \\item The range of observer and obliquity angles allowing for the largest possible pulsed fraction is determined by the angular width of the beamed component. In the case of a very narrow beamed component, the maximum pulsed fractions will be observed when the obliquity and observer angles are very close to each other. \\end{enumerate} Recently, Tauris \\& Manchester (1998) have reanalyzed radio pulsar polarization data to compute the obliquity distribution of the parent population of all radio pulsars. Their derived distribution peaks at small obliquities and suggests that most pulsars have $\\alpha \\lsim 35^0$. Our results are thus consistent with this picture. There are several important issues that we have thus far not addressed in our modeling of X-ray and $\\gamma$-ray pulse profles: polar cap heating and plausible physical reasons for the observed phase offset between X-ray and $\\gamma-$ray pulses. Since the pioneering study by \\cite{rud75} it has been understood that the development of electron-positron cascades above the NS surface initiated by primary electrons should unavoidably result in precipitation of ultra-relativistic positrons on the stellar surface. The kinetic energy of these positrons should be eventually transformed into the thermal energy of the NS crust and then reradiated most likely in soft X-rays (see e.g. Arons \\& Scharlemann 1979, and Arons 1981). The polar cap heating could thus add a component to the thermal X-ray pulse profiles. The efficiency of the polar cap heating depends on the number density of positrons that flow back from the pair formation front (PFF) to the stellar surface. This number density cannot be calculated from first principles. Instead it is rather sensitively determined by the transverse and longitudinal structure of the PFF, electrodynamics of the PFF, and by the dynamics of positron acceleration (Harding \\& Muslimov 1997b). Estimates of X-ray luminosity due to polar cap heating by Arons (1981) predict that such heating accounts for only $8\\%$, $0.12\\%$ and $0.005\\%$ of the observed thermal X-ray luminosity of Vela, Geminga and PSR 0656+14, respectively (Harding 1995). However, these polar cap heating estimates need to be revised using more recent calculations of the electric fields above the polar cap (Muslimov \\& Tsygan 1990, 1992, and Muslimov \\& Harding 1997). We can make the following rough estimates based on the results of the general relativistic treatment of the acceleration of the primary beam over the polar cap. The main difference between this and the classical treatment (see also Arons 1996) is that the general-relativistic dragging of inertial frames allows very efficient acceleration even for an aligned rotator and does not require the concept of ``favorably curved field lines\" introduced by \\cite{aro79}. In the regime of the space-charge-limitation of current, the electric field in the region of electron acceleration is determined by the difference between the local charge density of electrons and Goldreich-Julian charge density $\\rm \\Delta \\rho _e$. This difference reaches a maximum $\\rm |\\Delta \\rho _e|_{max} \\sim |\\rm \\rho _{GJ}| \\kappa $ (where $\\rm \\kappa \\approx 0.38 r_g/R \\sim 0.15$, $\\rm r_g$ is the stellar gravitational radius, see e.g. Muslimov \\& Harding 1997) at some height $\\rm h_m$ above the surface and then exponentially declines toward the PFF. The backflowing positrons enter a regime of relativistic motion near $\\rm h_m$, where $\\rm |\\Delta \\rho _e| \\sim |\\Delta \\rho _e|_{max}$. The backflowing positrons tend to reduce the difference $\\rm |\\Delta \\rho _e|$ and therefore the maximum value of the electric field. Thus, the maximum charge density of backflowing positrons can be estimated as $\\rm |\\Delta \\rho _p| \\sim |\\Delta \\rho _e|_{max}$. The total power put in the acceleration of these positrons can be estimated as (see also Muslimov \\& Harding 1997) $\\rm L^{+}\\sim f~L^{-}_{max}$, where $\\rm L^{-}_{max}$ is the maximum power of primaries, and $\\rm f \\approx 0.25 \\kappa $. Since $\\rm L^{-}_{max}\\sim L_{\\gamma}$, we get that $\\rm L^{+}\\sim f~L_{\\gamma}$. For a canonical NS of mass 1.4 $\\rm M_{\\odot }$ and radius 10 km we get the values $\\rm L^{+}$ that may account in part for the observed soft X-ray fluxes from pulsars (see e.g. Becker \\& Tr\\\"umper 1997). We now discuss the possible explanation for the phase offset (by about 0.15 in phase) of the X-ray pulse center toward the trailing $\\gamma$-ray peak, observed in Geminga and Vela (see \\S~2). In the framework of the polar cap model, the phase offset might be explained in terms of e.g. 1) an offset dipole 2) asymmetric polar cap heating and/or 3) dragging of photon geodesics by the gravitational field of a rotating body (see e.g. Misner, Thorne, \\& Wheeler 1973). For the offset dipole, the geometrical vertex of the $\\gamma-$ray emitting cone projects onto the stellar surface at a point offset from the magnetic pole. This means that the phase of the thermal X-ray emission (centered at the magnetic pole) will be offset from the $\\gamma-$ray pulses. A significant polar cap heating component may add to the thermal X-ray profile caused by cooling. If the heating by precipitating particles is not uniform over the polar cap, then the resulting X-ray component may add more at the trailing edge of the thermal pulse. We have not yet made a detailed calculation of the distribution of heating rates over the polar cap, so this effect is hard to predict at present. Any offset due to asymmetric particle heating might be enhanced by an asymmetric $\\gamma-$ray precipitation on the stellar surface near the polar cap from downward cascades due to positron acceleration. The effect of frame dragging on the light rays results in a phase shift of both X-ray and $\\gamma-$ray pulses. However, since the magnitude of this effect has not been accurately evaluated and is beyond the scope of this paper, we cannot say whether this effect will quantitatively account for the observed phase offset for a 0.1-0.3 s pulsar. In forthcoming publications we plan to discuss these and other effects in more detail and present more comprehensive theoretical analysis of the observed X- and $\\gamma $-ray emission from Geminga, Vela, PSR 0656+14, and other pulsars from this subpopulation of radio pulsars." + }, + "9802/astro-ph9802110_arXiv.txt": { + "abstract": "Deuterium abundance re-measurements by Burles \\& Tytler (1998; hereafter BT) yielded D/H = (3.3$ \\pm 0.3)\\times10^{-5}$ and the robust upper limit D/H $ < 3.9\\times10^{-5}$ from the $z_{\\rm a} = 3.572$ system toward Q~1937-1009. In this new analysis BT adopted multicomponent {\\it microturbulent} models together with the possibility to vary freely the local continuum level around each \\ion{H}{1} line to improve the fit. The procedure failed, however, to fit adequately D~Ly-$\\beta$ without recourse to an additional H~Ly-$\\alpha$ contamination at the position of D~Ly-$\\beta$. We show that this obstacle may be successfully overcome within the framework of the {\\it mesoturbulent} model accounting (in contrast to the microturbulent approximation) for a correlated structure of the large scale velocity field. Using the {\\it same} observational data and the original continuum as determined by Tytler et al. (1996), we obtained good fits. The one-component mesoturbulent models provide D/H in the range $\\simeq (3.2 - 4.8)\\times10^{-5}$ and the total hydrogen column density N(\\ion{H}{1}) $\\simeq (5.6 - 7.0)\\times10^{17}$ cm$^{-2}$. This result is consistent with that found by us from the $z_{\\rm a} = 2.504$ and $z_{\\rm a} = 0.701$ systems toward Q~1009+2956 and Q~1718+4807, respectively. The range for D/H common to all three analyses is D/H $\\simeq (4.1 - 4.6)\\times10^{-5}$. This value is consistent with standard big bang nucleosynthesis [SBBN] if the baryon-to-photon ratio, $\\eta$, is in the range $4.2\\times10^{-10} \\lesssim \\eta \\lesssim 4.6\\times10^{-10}$, implying $0.0155 \\lesssim \\Omega_{\\rm b}h^2_{100} \\lesssim 0.0167$. ", + "introduction": "The most distant absorption-line system with observable \\ion{D}{1} Ly-$\\alpha$ and Ly-$\\beta$ lines is the Lyman limit system discovered by Tytler et al. (1996, hereafter TFB) at $z_{\\rm a} = 3.572$ toward the quasar Q~1937-1009 ($z_{\\rm e} = 3.78$). In TFB, high resolution (FWHM = 9 km s$^{-1}$) and high signal-to-noise (S/N $\\simeq 75$ per pixel at the position of the Ly-$\\alpha$ line) Keck spectra of Q~1937-1009 revealed hydrogen absorption throughout the entire Lyman series as well as a few metal absorption lines with asymmetric profiles. Their first measurements of D/H based on a two-component microturbulent model (Voigt profile deconvolution analysis) gave a low value of D/H = $(2.3 \\pm 0.3 \\pm 0.3)\\times10^{-5}$ ($1\\sigma$ statistical and systematic errors). This result caused a lively discussion since the low D/H value would imply a high universal density of baryons $\\Omega_{\\rm b}h^2 = 0.024^{+0.006}_{-0.005}$ (TFB), where $\\Omega_{\\rm b}$ is the fraction of the critical density in form of baryons and $h$ is the Hubble parameter scaled to 100 km~s$^{-1}$~Mpc$^{-1}$. Note that the total SBBN baryon density compatible with the abundances of D and $^3$He observed in the solar system and the interstellar medium, has been estimated to be $\\Omega_{\\rm b}h^2 = 0.0175^{+0.010}_{-0.005}$ (Hata et al. 1996), which is slightly larger than earlier BBN estimates $\\Omega_{\\rm b}h^2 = 0.0125 \\pm 0.0025$ (Walker et al. 1991). The lack of uniqueness in the Voigt deconvolution procedure was employed by Wampler (1996) to show that other microturbulent models could give D/H ratios that are about 3 times higher than the TFB's value and still are compatible with their fit. Wampler also suggested that the total hydrogen column density may not be well determined by TFB because of either incorrect sky subtraction or improper modeling the Ly-$\\alpha$ forest structure above and below the Lyman continuum break. The latter is problematical for distant QSOs exhibiting a high density Ly-$\\alpha$ forest, and different methods of analysis have yielded by now different values for N(\\ion{H}{1})~: $(3.8 - 4.9)\\times10^{17}$ cm$^{-2}$ by Songaila et al. (1997); $(6.9 - 7.6)\\times10^{17}$ cm$^{-2}$ by Burles \\& Tytler (1997). A value of N(\\ion{H}{1}) $\\sim 6\\times10^{17}$ cm$^{-2}$ which ``produces a smooth forest opacity above and below the break'' was recently presented by Songaila (1998). Using new constraints on N(\\ion{H}{1})$_{\\rm tot}$, BT re-considered the D/H measurements from the $z_{\\rm a} = 3.572$ system. In this new approach, the metal absorption lines are not used to constrain D/H. The absorbers are separated into two groups -- with low and high \\ion{H}{1} column densities [N(\\ion{H}{1}) $\\leq 2.5\\times10^{15}$ cm$^{-2}$ and N(\\ion{H}{1}) $\\gtrsim 10^{17}$ cm$^{-2}$, respectively]. The first group does not show \\ion{D}{1} and was utilized to fit the Ly-$\\alpha$ forest features in the vicinity of the Lyman series lines from the $z_{\\rm a} = 3.572$ system. The second (main) group of lines were described by five free physical parameters~: N(\\ion{H}{1}), N(\\ion{D}{1}), $z$, $b_{\\rm tur}$, and $T$. The best fitting model with 3 main components (Model~4) allows for additional free parameters characterizing the local continuum to improve the fit. BT conclude, however, that ``the model fit to Ly$\\beta$ is not as good as Ly$\\alpha$, there is under-absorption in two places near D-Ly$\\beta$''. The present Letter is primarily aimed at the inverse problem in the analysis of the H+D Ly-$\\alpha$ and Ly-$\\beta$ absorption observed by TFB. It is shown that the mesoturbulent model based on only 5 physical parameters and an appropriate velocity field configuration is a sufficient description of the D~Ly-$\\alpha$ and Ly-$\\beta$ lines with no free parameters in the continuum. ", + "conclusions": "We have shown that the interpretation of the Q~1937-1009 spectrum obtained by TFB is not unique. The data can be modeled with a higher D/H ratio if one accounts for spatial correlations in the large scale velocity field. The most accurate and robust RMC solution (model {\\bf d}) was obtained with a total hydrogen column density outside the range found by Burles \\& Tytler (1997) but in good agreement with the Songaila (1998) estimate. Of course, our analysis gives a certain {\\it range} for N(\\ion{H}{1})$_{\\rm tot}$ and D/H (see Table~1). However, this range narrows if we require the D/H value to be compatible with the results of our previous analyses of the \\ion{D}{1} absorption toward Q~1009+2956 ($z_{\\rm a} = 2.504$) (Paper~III) and toward Q~1718+4807 ($z_{\\rm a} = 0.701$) (Levshakov et al., 1998). The range common to all three analyses is D/H $\\simeq (4.1 - 4.6)\\times10^{-5}$. This implies for Q~1937-1009 N(\\ion{H}{1}) $\\sim (5.8 - 6.5)\\times10^{17}$ cm$^{-2}$, showing that model {\\bf d} is well within this range. From SBBN it follows that D/H $\\simeq (4.1 - 4.6)\\times10^{-5}$ implies for the baryon-to-photon ratio, $\\eta$, a value in the interval $\\simeq (4.2 - 4.6)\\times10^{-10}$. With the present-day photon density determined from the cosmic microwave background (e.g. Fixen et al. 1996), one then estimates the present-day baryon density to be in the range $\\Omega_{\\rm b}h^2 \\simeq 0.0155 - 0.0167$. The final conclusion is that the current observations support SBBN and that there is no conflict with the D/H measurements within the generalized mesoturbulent approach." + }, + "9802/astro-ph9802056_arXiv.txt": { + "abstract": "We contrast measurements of composite optical and ultraviolet (UV) spectra constructed from samples of QSOs defined by their soft X-ray brightness. X-ray bright (XB) composites show stronger emission lines in general, but particularly from the narrow line region. The difference in the \\oiii/\\hb\\, ratio is particularly striking, and even more so when blended \\feii\\, emission is properly subtracted. The correlation of this ratio with X-ray brightness were principal components of QSO spectral diversity found by Boroson \\& Green. We find here that other, much weaker narrow optical forbidden lines (\\oii\\, and \\nev) are enhanced by factors of 2 to 3 in our XB composites, and that narrow line emission is also strongly enhanced in the XB UV composite. Broad permitted line fluxes are slightly larger for all XB spectra, but the narrow/broad line ratio stays similar or increases strongly with X-ray brightness for all strong permitted lines except \\hb. Spectral differences between samples divided by their relative X-ray brightness (as measured by \\aox) exceed those seen between complementary samples divided by luminosity or radio loudness. We propose that the Baldwin effect may be a secondary correlation to the primary relationship between \\aox\\, and emission line equivalent width. We conclude that either 1) \\ew\\, depends primarily on the {\\em shape} of the ionizing continuum, as crudely characterized here by \\aox\\, or 2) both \\ew\\, and \\aox\\, are related to some third parameter characterizing the QSO physics. One such possibility is intrinsic warm absorption; a soft X-ray absorber situated between the broad and narrow line regions can successfully account for many of the properties observed. ", + "introduction": "\\label{intro} Most QSOs have been discovered by virtue of their strong optical/UV emission lines, or non-stellar colors in this bandpass. Our understanding to date of the violent inner regions of active galactic nuclei (AGN) also derive in large part from their optical/ultraviolet (OUV) spectra. The production of QSO emission lines is widely attributed to photoionization and heating of the emitting gas by the UV to X-ray continuum (e.g., Ferland \\& Shields 1985; Krolik \\& Kallman 1988). Studies investigating the relationship of emission lines to continuum radiation have a long history in the field, but several strong observational relationships remain unexplained. If the proportionality between line and continuum strength were linear, then diagnostics such as line ratios and \\ew\\, would be independent of continuum luminosity. Baldwin (1977) first noticed that in high redshift quasars, the \\ew\\, of the CIV $\\lambda1550$\\AA\\, emission line in quasars decreases with increasing UV ($1450$\\AA) luminosity. The Baldwin effect (BEff, hereafter) was also found to be strong for ions such as OVI, NV, He\\,II, CIII], Mg\\,II, and Ly$\\alpha$ (e.g., Tytler \\& Fan 1992, Zamorani et al. 1992). The initial excitement about the potential for the BEff as a standard candle and cosmological probe has faded; the dispersion in the relationship is too large. Neither the source of that dispersion, nor cause of the BEff itself have been definitively identified. However, some possible explanations for the BEff have been offered, one being a dependence of the continuum spectral energy distribution (SED) on luminosity (Zheng \\& Malkan 1993, Green 1996). Many important lines respond primarily to the extreme ultraviolet (EUV) or soft X-ray continuum. Unfortunately, the EUV band is severely obscured by Galactic absorption. However, constraints on the ionizing continuum are available through analysis of the adjacent UV and soft X-ray windows. In a small, uniform sample of optically-selected QSOs (Laor et al. 1997), the strongest correlation found between X-ray continuum and optical emission line parameters was of the soft X-ray spectral slope \\ax\\, (where $\\fnu \\propto \\nu^{-\\ax}$) and the FWHM of the \\hb\\, emission line. Strong correlations between \\ax, $L_{\\oiii}$, \\feii/\\hb, and the \\oiii/\\hb\\, ratio were seen both there and in previous QSO studies (e.g., Boroson \\& Green 1992). The latter authors found that most of the variation in the observed properties of low-redshift QSOs can be represented in a principal component analysis by eigenvectors linking \\feii, \\oiii, \\hb, and \\heii\\, emission line properties and continuum properties such as radio loudness, and the relative strength of X-ray emission, as characterized by \\aox\\, (defined below). A recent, and possibly related result is that Seyfert 1s with broad \\hb\\, emission lines tend to have hard (flat) X-ray spectral slopes (e.g., Brandt, Mathur, \\& Elvis 1997). In a hard-X-ray-selected sample of (mostly Seyfert) AGN, narrow \\oiii\\, flux correlates well with X-ray flux, while broad Balmer lines do not (Grossan 1992). The physical origin of these diverse and interrelated correlations has yet to be determined. We are launching a large-scale effort to probe these effects in large samples, using both data and analysis as homogeneous as possible. Many physically informative trends intrinsic to QSOs may be masked by dispersion in the data due to either low signal-to-noise or variability. An important tool for studying global properties of QSOs is the co-addition of data for samples of QSOs. In this paper, we concentrate on an analysis of composite optical/UV spectra of subsamples of QSOs grouped by the relative strength of their soft X-ray emission. ", + "conclusions": "\\label{discuss} We now consider a variety of possibilities to account for the optical and UV spectral difference between X-ray bright and X-ray faint QSOs; 1) luminosity effects, 2) radio loudness, 3) absorption, and 4) changes in the {\\em intrinsic} spectral energy distribution (SED). The strength of some of these effects is directly testable using the samples at hand. \\subsection{Luminosity Subsamples} \\label{lum} Both selection effects and secondary correlations must be considered when evaluating the significance of observed correlations such as these. Two well-known effects could conspire to produce an overall weakening of emission lines with increasing \\aox. First, \\aox\\, is known to increase with optical luminosity (Wilkes et al. 1994, Green et al. 1995), at least for optically-selected samples (LaFranca 1995). Secondly, as luminosity increases, line equivalent width decreases (i.e., the Baldwin Effect; Baldwin 1977). Could these effects combine to produce the anti-correlation of \\aox\\, and line strength observed here? As can be seen from Table~1, the LBQS subsamples are well-matched in optical luminosity, so that no Baldwin effect is expected. Furthermore, the strength of the Baldwin effect in the optical is known to be weak. Our results from the \\IUE\\, subsamples are less immune to a Baldwin effect/SED conspiracy for the following reasons: 1) our \\IUE\\, subsamples are not as well-matched in luminosity, and 2) the Baldwin effect is much stronger in the UV. We therefore perform a stringent check, by applying the same spectral averaging techniques now to new subsamples defined by UV luminosity. We divided the \\IUE\\, sample at the mean UV luminosity \\footnote{The log rest-frame luminosity \\logluv, in \\fnucgs\\, at 1450\\AA, is defined in Green (1996)} value of 30.6. The resulting low UV luminosity (UVLO; 27 QSO) and high UV luminosity (UVHI; 22 QSO) subsamples both had mean $\\aox=1.4\\pm .04$. SPECFIT procedures identical to those of the XB and XF samples were applied. Virtually all spectral differences were {\\em less} significant in the UV luminosity subsamples than in the \\aox\\, subsamples. Only narrow \\lya\\, emission changes more strongly between UV subsamples than between the \\aox\\, subsamples. Indeed, the Baldwin Effect appears to be strongest in the narrow line components of both \\lya\\, and \\civ. The bulk of the effect could be due to differences in narrow line region (NLR) emission, as also suggested by Osmer, Porter, \\& Green (1994). Since emission line correlations are stronger with \\aox\\, than with luminosity, we conclude that either 1) \\ew\\, depends primarily on the {\\em shape} of the ionizing continuum, crudely characterized here by \\aox\\, or 2) both \\ew\\, and \\aox\\, are related to some third parameter characterizing the QSO physics. One such possibility is absorption. Could the correlation of \\aox\\, to luminosity {\\em cause} the Baldwin effect? Although by design we have selected subsamples of similar luminosity for our XF and XB composites, we may suppose that the primary relationship between \\ew\\, and \\aox\\, is propagated into the relationship between \\aox\\, and luminosity, and test the strength of the predicted secondary relationship between \\ew\\, and luminosity that results, i.e. the Baldwin Effect. We begin by simply contrasting the observed response here of \\wciv\\, and \\wmgii\\, to \\aox\\, with that predicted in the most recent comprehensive study of the Baldwin effect in optically-selected QSOs (Zamorani et al. 1992). The change $\\Delta\\overline{\\aox}\\sim 0.3$ between XF and XB subsamples is similar in our \\IUE\\, and LBQS samples. The change in log\\,\\wmgii\\, predicted by the BEff in this line is $\\Delta{\\rm log}~\\wmgii) \\sim -0.68\\aox = 0.18$. The change we actually measure between XB and XF composites is (from the values in Table~3) $\\Delta{\\rm log}~\\wmgii=1.0$. Similarly, the change predicted by the BEff in log\\,\\wciv\\, seen by Zamorani et al. is $\\Delta{\\rm log}~\\wciv \\sim -1.16\\aox = 0.30$, while we actually measure $\\Delta{\\rm log}~\\wciv=1.57$. Thus, the effect of \\aox\\, on emission line strengths is some 5 -- 6 times larger than that expected if it were secondary to a luminosity effect. Indeed, {\\em it seems likely that the BEff is secondary to the relationship between \\aox\\, and line equivalent width.} If the BEff is caused by a change in \\aox\\, concomitant with luminosity, then the strongest \\ew(\\aox) relationship we see, that of $\\woiii$, predicts a BEff in \\oiii. This is indeed seen in the BG92 data, where the probability of no correlation (the null hypothesis) between $M_V$ and \\woiii\\,$\\lambda 5007$ is $<1\\%$. \\subsection{Radio Properties} \\label{radio} Radio and X-ray loudness are correlated; Green et al. (1995) confirmed that RL QSOs are more soft X-ray luminous than RQ QSOs in the LBQS. Unfortunately, the difference in radio loudness between the LBQS XF and XB subsamples cannot be well-characterized: although 4 (of 17 with radio data) are radio-detected in the LBQS XB subsample, only one QSO (of 38) in the XF subsample is a radio detection. However, OUV spectral differences as a function of radio loudness have been extensively studied. Differences between emission lines in radio loud QSO (RLQ) and radio quiet QSO (RQQ) spectra longward of 1600\\AA\\, are quite subtle, and there is a remarkable similarity in \\mgii\\, and \\ciii\\, emission lines between RLQs and RQQs (Francis et al. 1993). Distinctions found by BG92 include a redward asymmetry of \\hb\\, in RLQs, whereas RQQs show about equal numbers of red and blue asymmetries (BG92). RL QSOs tend to have strong \\oiii\\, and weak optical \\feii. In Wang, Brinkmann, \\& Bergeron (1996), neither \\hb\\, nor optical \\feii\\, measurements correlate with radio loudness. To check the effects of radio loudness on UV emission lines, we created RL and RQ composite UV spectra from the \\IUE/\\ein\\, sample. We determined radio loudness from Falcke et al. (1996) for all the UV excess selected (PG) QSOs, and from Veron-Cetty \\& Veron (1993) or the NED database for others. Two QSOs are omitted because of intermediate radio loudness, and two more have no published radio data. Radio subsamples are identical in luminosity and well-matched in redshift if we impose $z>0.15$. This yields N=26 and $\\overline{z}=0.48$ for the RL subsample, N=13 and $\\overline{z}=0.52$ for RQ QSOs, and \\loglopt=30.9 for both. There is some remaining difference between \\aox\\, distributions for the \\IUE/\\ein\\, radio subsamples, as expected, such that the RL sample is somewhat more X-ray bright ($\\overline{\\aox}=1.36\\pm0.04$) than the RQ sample ($\\overline{\\aox}=1.47\\pm0.05$). We find \\lya\\,emission to be somewhat stronger in the RQ composite. \\civ\\, emission is of similar strength and FWHM in both RQ and RL composites, but displaced slightly to the red in the RL composite. Similar results were found independently by Wills \\& Brotherton (1995). Thus overall, the emission line trends due to radio loudness in our sample are weaker than, and tend to {\\em diminish} those seen in \\aox. The important result is that {\\em spectral differences between composites binned by \\aox\\, are significantly larger than between composites binned by radio loudness}. \\subsection{Absorption and the \\aox(\\lopt) Relation} Absorption by ionized gas near the nucleus can extinguish soft X-ray emission without having significant effects on the observed optical emission. The BAL QSOs represent an extreme example, wherein observed soft X-ray fluxes are at least an order of magnitude below those of non-BAL QSOs of similar optical brightness (Green \\& Mathur 1996). Even AGN with much less spectacular UV absorption (e.g. narrow absorption lines) show significant soft X-ray absorption (Mathur 1994, Mathur, Elvis, \\& Wilkes 1995). Since most QSOs have either low S/N, low resolution, or no available UV spectra, many such absorbers await recognition in current QSO samples. Is it possible that the \\aox(\\lopt) correlation is itself caused by absorption? The hypothesis might be tested if {\\em all} types of absorbed QSOs were removed, including QSOs with known narrow-line intrinsic, damped \\lya, and \\lya\\, forest absorbers. A higher S/N sample that has been systematically searched for absorption in both UV and soft X-ray spectra is needed. We are building such a sample from the HST FOS and ROSAT public archives, and will study these issues in an upcoming paper. We note that if at least some of the correlation is caused by increased warm absorption at higher luminosities, then the correlation would be expected to be further weakened in soft X-ray selected samples. More common absorption in the UV bandpass is suspected anecdotally at higher luminosities, but remains to be demonstrated statistically. In the X-ray bandpass, the same can be said for the popular assumption that absorption is {\\em less} common at high luminosity. This latter notion should be particularly suspect, since higher energy rest-frame X-ray emission is observed in most high luminosity objects, and requires a higher intrinsic absorbing column for detection. Furthermore, the most luminous objects in samples to date have also been radio-loud (Reynolds 1997). Variability may play a significant role in the \\aox(\\lopt) correlation. Typically, optically-selected AGN are found near the survey flux limit and therefore preferentially in a bright phase, and then later followed up in an X-ray pointed observation. Overall, more distant (and thus more luminous) optically-selected QSOs would tend to have larger measured \\aox. Even if variability-related biases are not responsible for the observed \\aox(\\lopt) correlation, such variability may be expected to introduce scatter into the true relation. As such, $\\alpha_{ix}$ (the powerlaw spectral index between 1$\\mu$m and 2~keV) might prove a more variability-resistant measure. Indeed, Lawrence et al. (1997) have found that some of the primary correlations discussed here become more significant when $\\alpha_{ix}$ replaces \\aox. Not only variability, but also absorption more strongly affects the optical than the infrared. \\subsection{Intrinsic Spectral Energy Distributions} \\label{seds} We have offered one interpretation of our measurements that assumes that the {\\em intrinsic} broadband continuum emission from the QSO central engine is constant in shape, independent of luminosity, but that the spectral energy distribution (SED) seen by the NLR and/or by us may be strongly affected by intervening, possibly ionized absorbing clouds. However, the change in \\aox, and the accompanying changes in emission lines may at least in part be due to changes in the intrinsic SED. Many QSOs have soft ($\\lapprox$1~keV) X-ray emission that exceeds the extrapolation from the powerlaw continuum observed at higher energies (e.g., Turner \\& Pounds 1989, Masnou et al. 1992). This X-ray `soft excess' has often been interpreted as the high energy continuation of the UV/EUV/soft X-ray ``big blue bump'' (BBB), possibly thermal emission from the surface of an accretion disk (although see Barvainis 1993). Several workers (beginning with Malkan \\& Sargent 1982) have proposed as an explanation of the Baldwin effect that as QSO luminosity increases, the BBB shifts toward lower energies at higher (OUV) luminosities. As QSOs become more luminous in the optical/UV band, they thereby undergo a weaker increase in \\lx, and therefore an increase in \\aox. The response of line flux and \\ew\\, depends in a fairly complicated manner on the peak energy of the BBB, on the BBB normalization relative to the powerlaw continuum, and on the ionization and heating continuum of the line species in question. Detailed photoionization modeling using a variety of input continua, impinging on an ensemble of clouds from the broad to the narrow line region, including full self-shielding and optical depth effects are called for (see e.g., Korista et al. 1997, Baldwin et al. 1995, Shields et al. 1995). On average, the strongest effect may be that higher luminosity QSOs may undergo spectral evolution such that fewer photons from a soft X-ray excess/BBB component are available for ionization. A shift of the BBB to lower energies at higher luminosity implies that the soft X-ray excess should decrease to higher luminosities. An apparent hardening of soft X-ray PL spectral index has been seen in composite ROSAT spectra of LBQS QSOs toward higher redshifts (e.g., Schartel et al. 1996) that gibes with this picture, but again, higher energy rest-frame X-ray emission is observed in most high luminosity objects, where a higher intrinsic column is required before absorption can be detected. In this picture, a stronger BEff might be expected for species of higher ionization energy. There is some evidence for such a trend (Zheng et al. 1992). The intrinsic SED model does not predict absorption features, but does imply that the \\aox(\\lopt) relation should persist even in soft X-ray-selected samples." + }, + "9802/astro-ph9802260_arXiv.txt": { + "abstract": "New bolometer arrays operating on the world's largest sub-millimetre and millimetre telescopes offer a unique view of the high-redshift universe with unprecedented sensitivity. Recent sub-millimetre continuum studies show that the host galaxies of many luminous high-redshift active galactic nuclei (radio galaxies and radio-quiet quasars) radiate strongly at rest-frame far-infrared wavelengths and thus contain substantial quantities of dust. In the majority of these high-redshift AGN-hosts, the inferred star formation is proceeding at a rate comparable to that found in local, interacting ultra-luminous far-infrared galaxies. This level of activity is an order of magnitude greater than the more modest star-formation rates apparently displayed by the recently-discovered Lyman-limit galaxies at ${\\rm z} \\sim 3$, which have been argued to represent the era of spheroid formation (although the degree to which the effects of reddening by dust grains may have biased the interpretation of these optical/UV studies of high-z galaxies has yet to be properly determined). However, it is too early to say whether such bright far-infrared emission is a feature of all massive galaxies at $z > 3$, or whether it is in fact confined to the hosts of the most luminous AGN. In this paper we review the current status of cosmological observations at sub-millimetre and millimetre wavelengths, highlighting our own recent SCUBA observations of high-redshift radio galaxies. We also explain how observations over the next few years should allow the true level of star-formation activity in the high-redshift universe to be properly quantified, and we provide example predictions for the first deep sub-millimetre survey (of the Hubble Deep Field), which we and our colleagues are currently undertaking with SCUBA at the JCMT. ", + "introduction": "Sub-millimetre (sub-mm) cosmology is still in its infancy. Despite significant efforts with single-element bolometers over the last 5 years, only a handful of unlensed objects have been unambiguously detected (- for a summary, see Hughes, Dunlop \\& Rawlings 1997). However, the recent advent of sensitive sub-mm/mm bolometer arrays seems set to revolutionize the field. Studies of statistically significant samples of known high-redshift sources are now feasible, as are the first meaningful sub-mm blank-field surveys. The key logical steps in our own approach to this burgeoning field can be summarized as follows: \\begin{itemize} \\item Make pointed observations of known high-redshift objects spanning a wide range in redshift. We have concentrated on high-redshift steep-spectrum radio galaxies primarily because they can be reasonably expected to be the progenitors of at least a subset of present-day massive ellipticals, but also because, being selected on the basis of extended emission, their sub-mm properties are unlikely to be biased by gravitational lensing. \\item Use sub-mm/far-infrared(FIR) data to first identify the emission mechanism that dominates the production of the rest-frame sub-mm/FIR luminosity. Then, if this emission is proven to be optically-thin thermal re-radiation from dust grains, use these same data to constrain the dust temperature, and hence calculate a reddening-free measure of the dust mass and infer the total mass of molecular gas in the galaxy. Assuming the dust grains are heated by young, massive stars, and not by an AGN, then estimate the `current' star-formation rate (SFR). \\item Compare the gas masses available for further star formation in the host galaxies of high-redshift active galactic nuclei (AGN), with the the final stellar masses of their expected low-redshift counterparts. Hence infer the evolutionary status of the galaxies that host high-redshift AGN. \\item Building on the results of such pointed observations, design and undertake a series of complementary sub-mm surveys reaching different depths over different areas. Use the measured source counts and redshift distributions to determine the true level and history of star-formation activity at high redshift, and to determine the form of the cosmological evolution of the dust-enshrouded starburst population. \\end{itemize} ", + "conclusions": "" + }, + "9802/astro-ph9802307_arXiv.txt": { + "abstract": "Recently Schlegel, Finkbeiner \\& Davis published an all-sky reddening map based on the COBE/DIRBE and IRAS/ISSA infrared sky surveys. Using the reddening map of Baade's Window and sample of 19 low-latitude ($|b|<5\\deg$) Galactic globular clusters I find that the DIRBE/IRAS reddening map overestimates $E(B-V)$ at low galactic latitudes by a factor of $\\sim 1.35$. I also demonstrate the usefulness of this high resolution map for selecting low-reddening windows near the Galactic plane. ", + "introduction": "We live in a dusty Universe (Hoover 1998, private communication), and correcting for the dust extinction and reddening affects almost all aspects of optical astronomy. For us, observing from within the Milky Way, it is of crucial importance to know how much Galactic dust there is towards various objects. Burstein \\& Hailes (1982; hereafter: BH) constructed an all-sky reddening map, used extensively by the astronomical community.\\footnote{Their paper was cited 540 times between 1992 and 1997} Recently, Schlegel, Finkbeiner \\& Davis (1998; hereafter: SFD) published a new all-sky reddening map, based on the COBE/DIRBE and IRAS/ISSA maps.\\footnote{The reddening map and related files and programs are available using the {\\tt WWW} at: {\\tt http://astro.berkeley.edu/davis/dust/}} This map is intended to supersede the BH map in both the accuracy and the spatial resolution ($6.1'$). Also, unlike the BH map, which presented the $E(B-V)$ values only for the $|b|>10\\deg$ regions, the SFD map extends all the way to the Galactic plane. There are many instances when one would like to be able to identify low-reddening regions near the Galactic plane. As an example, the expected microlensing optical depth increases strongly towards the Galactic center (e.g. Stanek et al.~1997, their Figure~12). In this paper I show the usefulness of the SFD map for selecting low-reddening regions close to the Galactic plane. In Section~2 I compare the predicted and the observed reddenings and arrive at the approximate scaling relation. In Section~3 I show the re-scaled reddening map for central region of our Galaxy and discuss its applications for microlensing experiments. ", + "conclusions": "" + }, + "9802/astro-ph9802131_arXiv.txt": { + "abstract": "The use of relations between structural parameters of early type galaxies to perform the Tolman test for the expansion of the Universe is reconsidered. Scaling relations such as the Fundamental Plane or the Kormendy relation, require the transformation from angular to metric sizes, to compare the relation at different z values. This transformation depends on the assumed world model: galaxies of a given angular size, at a given z, are larger (in kpc) in a non-expanding universe than in an expanding one. Furthermore, the luminosities of galaxies are expected to evolve with z in an expanding model. These effects are shown to conspire to reduce the difference between the predicted Surface Brightness change with redshift in the expanding and non expanding cases. We have considered expanding models with passive luminosity evolution. We find that their predictions for the visible photometric bands are very similar to those of the static model till z$\\sim$1, and therefore, the test cannot distinguish between the two world models. Recent good quality data on the Kormendy relation and the Fundamental Plane at intermediate redshifts are consistent with the predictions from both models. In the K-band, where the expected (model) luminosity evolutionary corrections are smaller, the differences between the expanding and static models amount to $\\sim$0.4 (0.8) magnitudes at z = 0.4 (1). It is shown that, due to that small difference between the predictions in the covered z-range, and to the paucity and uncertainties of the relevant SB photometry, the existing K-band data is not adequate to distinguish between the different world metrics, and cannot be yet used to discard the static case. It is pointed out that the scaling relations could still be used to rule out the non-evolving case if it could be shown that the coefficients change with the redshift. ", + "introduction": "The Big Bang cosmological model rests upon the Cosmological Principle, which determines the form of the metric of the world, and allows the explanation of the redshift phenomenon as a pure geometrical effect directly related to the nature of that metric. The Microwave Background Radiation or the abundance of the light elements strongly support the standard model. Thus, given the overall consistency of the model, some authors consider as superfluous any attempt to test the basic hypothesis, i.e., the form of the metric. Furthermore, its explanatory capacity is often used to rule out any other theoretical alternative. Recently, van Dokkum and Franx (1996) have argued that the microwave background radiation would not be planckian if the Surface Brightness (SB) dimming were different of $(1+z)^{-4}$. We remark that this is true from within the standard model, otherwise it should have been abandoned long ago. But the fact that data can be accommodated within the standard model, important as it is, cannot be considered as a formal proof of the underlying hypothesis. It is worth to remember that, formally, a false statement can have consequences that are true. Direct tests have been proposed since long (see Hubble \\& Tolman 1935), and recently reformulated from different angles (Moles 1991; Sandage \\& Perelmuter 1991; Kj\\ae rgaard, J\\o rgensen \\& Moles 1993). Among the tests proposed by Hubble \\& Tolman, the SB test has been considered the most powerful and {\\it clean}. Indeed, the SB of a standard candle only depends on the red-shift: it would decrease as $(1+z)^{-4}$ in an expanding Universe. Otherwise, i.e., in a static, non-expanding universe, with some (unknown) operating mechanism to produce the observed redshift, the SB dimming would go like $(1+z)^{-1}$. The test cannot be considered as able to select one among different models since there are no known alternatives to the standard cosmology based on a non-expanding metric. Here, as stated by Hubble \\& Tolman, the non-expanding case is taken to gauge the standard model. In practice, when the Fundamental Plane (FP) or Kormendy relations are used, it is necessary to use the metric to transform the angles we measure into linear sizes, to compare data at different z values. Thus, what we actually have is a mixture of the SB and angular size tests. This is discussed in the next section, where the test is formulated in terms of some relation (assumed universal) between structural parameters of early type galaxies. The evolutionary aspects and corrections are discussed in \\S 3. The results, based on good quality data recently published (Pahre et al. 1996; van Dokkum and Franx 1996; Fasano et al. 1997; Kelson et al 1997; Bender et al 1997) are presented in \\S 4. A discussion of the results together with the conclusions can be found in \\S 5. Previous claims (Pahre et al 1996) that the static case is discarded at the 5$\\sigma$ confidence level are shown to be incorrect. Particular emphasis is made on the almost-lack of decidability of the test up to $z\\sim$1, except perhaps in the K-band. ", + "conclusions": "We have shown that the dependence on cosmology of the Kormendy relation or the FP includes, besides the SB cosmic dimming and the evolutionary (if any) and K-corrections, a factor depending on the metric. This metric term, that has been generally overlooked, has a sizeable effect, when the static predictions are also considered. It explains the discrepancy between the results reported by Pahre et al (1996) and ours, as they did not take it into account. In principle, the SB test could be formulated in such a way to avoid the presence of the metric term, by using the $\\mu_e$-$\\sigma$ relation. However, looking at this relation for the clusters considered here, it is found that it is too loosely defined, and cannot be used with any confidence for the Tolman test. Our main result is that as far as scaling relations are used, the Tolman test cannot be conclusive in any practicable range of redshift for visible photometric bands data, due to the conspiracy of the metric term and the evolution corrections to reduce the difference between the static and expanding world models. In the V- and R- bands those differences are nearly zero till z$\\sim$1. This is also the case for B-band data till z$\\sim$0.6. At larger redshifts the static case is enclosed in the spread of possible evolutionary corrections. The B-, V-, and R-band data are shown to fit very well the (similar) predictions of the different models. Only in the K-band the predictions are separated, even if only by modest amounts till z$\\sim$1, but on the empirical side, the situation is more ambiguous. The data are not fitted by any prediction, even if they are closer to the expanding case. We have noticed that there is a clear mismatch between the average $$ color index obtained from the 1kpc intercepts of the Kormendy relation, and the measured values for similar galaxies, or even the predictions at z$\\sim$0 by the evolutionary models. The situation needs to be clarified before a conclusion based on the Tolman test could be attempted. The data confirm that the FP and Kormendy relations are well defined at least till z$\\sim0.5$. In the static case they are automatically satisfied since the stationarity hypothesis implies that there should be always and in every cluster a population of galaxies satisfying a given relation. In the expanding case it represents only a constraint on the allowed evolutionary paths of early type galaxies, which should be checked independently of the Tolman test. It is clear that, if the change of the coefficients of the structural relation with z hinted by van Dokkum and Franx (1996) were confirmed, the static, non-evolving case would be ruled out. In that sense, the scaling relations constitute a test in themselves to prove the evolving character of the world. {\\sl Acknowledgments}. This research was partly supported by DGICYT (Spain) grant PB93-0139, and the Danish Natural Science Research Council grant 9401635. Carlos Barcel\\'o is acknowledged for helpful discussions on the nature of the tests. We acknowledge the referee for enlightening comments that did help to appreciably improve the manuscript. \\clearpage" + }, + "9802/astro-ph9802288_arXiv.txt": { + "abstract": "We discuss the consequences of almost rectilinear acceleration of protons to extremely high energies in a reconnection region on the surface of an accretion disk which surrounds a central black hole in an active galaxy. The protons produce $\\gamma$-rays and neutrinos in interactions with the disk radiation as considered in several previous papers. However, in this model the secondary $\\gamma$-rays can initiate cascades in the magnetic and radiation fields above the disk. We compute the spectra of $\\gamma$-rays and neutrinos emerging from regions close to the disk surface. Depending on the parameters of the reconnection regions, this model predicts the appearance of $\\gamma$-ray and neutrino flares if protons takes most of the energy from the reconnection region. In contrast, if leptons take most of the energy, they produce pure $\\gamma$-ray flares. The $\\gamma$-ray spectrum expected in the case of hadronic cascading is compared with the spectrum observed during the flare in June 1991 from 3C 279. The neutrino flares which should accompany these gamma-ray flares may be detected by future large scale neutrino telescopes sensitive at $\\sim 10^{5}$ TeV. ", + "introduction": "Flares on very short time scales in TeV $\\gamma$-rays have been recently detected from Mrk 421, Mrk 501, and 1ES2344+514 (e.g. Buckley et al.~1996, Gaidos et al.~1996, Aharonian et al.~ 1997, Catanese et al.~1998) and in the EGRET energy range from many other blazars (e.g. von Montigny et al.~1995, Mattox et al.~1997). Flares observed with a $\\sim 1$~day time scale are usually interpreted in terms of a relativistic shock moving through the jet with a Lorentz factor of the order of $\\sim 10$ (e.g. Maraschi, Ghisellini \\& Celloti~1992, Mannheim \\& Biermann~1992, Dermer \\& Schlickeiser~1993, Bednarek~1993, Sikora, Begelman \\& Rees~1994, Blandford \\& Levinson~1995, Ghisellini \\& Madau~1996, Protheroe~1997). However, the production of such flares may also be a natural consequence of energization of particles in a way similar to that causing solar flares, but in this case occurring in regions of magnetic reconnection which may exist on the disk surface (Haswell, Tajima \\& Sakai~1992, Lesch \\& Pohl~1992). Reconnection of magnetic field may also occur in the jet (e.g. Romanova \\& Lovelace~1992, Bisnovatyi-Kogan \\& Lovelace~1995). The processes of $\\gamma$-ray production by leptons accelerated in electric fields generated in the jet has been recently discussed by Bednarek \\& Kirk (1995, henceforth BK95) and Bednarek, Kirk \\& Mastichiadis (1996, henceforth BKM96). The acceleration of particles by almost rectilinear electric fields has some advantages in comparison to the stochastic acceleration by a relativistic shock since, in principle, particles can reach very high energies on shorter time scales. However, the process of reconnection of magnetic fields in astrophysical environments is poorly understood. Recently, Haswell, Tajima \\& Sakai (1992) have developed a model of explosive reconnection of magnetic fields on the surface of an accretion disk. They argued that particles (electrons, protons) can be accelerated to extremely high energies ($\\sim 10^{20-21}$ eV). The consequences of this model for $\\gamma$-ray production in blazars has been investigated recently by Bednarek (1997, henceforth B97), assuming that accelerated leptons develop cascades in the magnetic field through production of $\\gamma$-rays by the synchrotron process, and by the creation of secondary pairs by $\\gamma$-rays in the magnetic field. The $\\gamma$-ray spectra, emerging from a region on the inner part of an accretion disk around a massive black hole, are in agreement with observations of blazars. In this paper, we discuss the consequences of acceleration of protons in almost uniform electric fields induced by magnetic reconnection on the disk surface. Protons accelerated to extremely high energies produce $\\gamma$-rays and neutrinos through the decay of pions produced in collisions with disk photons. Neutrinos freely escape, but $\\gamma$-ray photons initiate cascades in the magnetic and radiation fields above the disk (B97). The conditions for acceleration of protons, and their propagation in the disk radiation, are discussed in Sect.~2.1 and 2.2. We compute the spectra of neutrinos and $\\gamma$-rays (Sect.~2.3), and describe the details of the cascades initiated by these $\\gamma$-rays (Sect.~3). Example $\\gamma$-ray spectra are compared with observations of the June 1991 flare from 3C 279, and a prediction of the neutrino emission accompanying this flare is given in Sect.~4. ", + "conclusions": "Some models of magnetized accretion disks predict the appearance of relatively small, but efficient, acceleration regions close to the surface of the inner part of the accretion disk as a result of magnetic reconnection (Haswell, Tajima \\& Sakai~1992, Lesch \\& Pohl~1992). In this paper, we discuss the consequences of acceleration of hadrons in such regions. Depending on the conditions in the acceleration region, and above the disk (see Fig.~\\ref{fig1}), accelerated hadrons may either escape, initiate cascades after emerging from the acceleration region, or produce copious secondary leptons during the acceleration process which initiate subsequent leptonic cascading. Primary leptons can also initiate cascades, and this possibility has been already discussed~(B97). In the analysis of hadronic cascades we included the processes of pion production by hadrons, synchrotron radiation of secondary $e^\\pm$ pairs in the quantum domain, magnetic $e^\\pm$ pair production by $\\gamma$-ray photons and absorption of the $\\gamma$-rays in the disk radiation. The $\\gamma$-ray spectra emerging from such a sequence of processes resemble the spectra observed from FSRQ blazars (Figs~5). Their comparison with the June 1991 flare from 3C 279 allows one to predict the strength of the accompanying neutrino flare (see Fig.~\\ref{fig6}). In the present model we have neglected the contribution of interactions of accelerated hadrons with the background matter, either during the acceleration stage, or after propagation through the disk radiation. Interactions of relativistic protons with the disk radiation dominate over the interactions with matter if the density of matter $n_{\\rm H}$ is lower than \\begin{eqnarray} n_{\\rm H}\\approx 2\\times 10^{11} T_4^3 ~{\\rm particles~ cm^{-3}}, \\end{eqnarray} \\noindent If the density of matter is relatively high at larger distances from the disk, then protons, at this stage having lost most of their energy during propagation through the disk radiation, can also produce neutrinos by interactions with matter. These neutrinos are typically of lower energies, and their power will be much less than in the higher energy neutrinos from interactions with radiation. \\begin{figure} \\vspace{7.cm} \\special{psfile=neutrino.eps voffset=0 hoffset=-20 hscale=60 vscale=60 angle=0} \\caption[]{The $\\gamma$-ray spectrum obtained from the present model (full curve) is compared with observations of the June 1991 flare from 3C 279 detected by COMPTEL (C) (Williams et al.~1995), EGRET (E) (Kniffen et al.~1993) and Ginga (G) (Hartman et al.~1996). The parameters of the fit are as in Fig.~\\ref{fig5}a. The dashed curve shows the expected muon neutrino flux during this flare. The dot-dashed curves show the atmospheric neutrino background (horizontal - upper curve, and vertical - lower curve) for a neutrino detector with $1^\\circ$ angular resolution .} \\label{fig6} \\end{figure} Here we have considered only a single acceleration region in order not to complicate the picture too much. However, in principle many acceleration regions with different parameters could be present simultaneously on the disk surface. Different types of cascades may then contribute to the observed $\\gamma$-ray spectrum from a specific blazar, and to the neutrino flux. If multiple reconnection sites do contribute, then the neutrino fluxes predicted for a single site should be considered as upper limits. However, a flare being produced by a single reconnection region seems to be equally probable given the likely short duration of each reconnection region ($\\sim l_{\\rm rec}/c$). It is expected that as a result of magnetic reconnection a large amount of energetic plasma may enter the jet. Therefore, the cascade processes taking place above the reconnection region should also occur in the jet plasma which moves with mildly relativistic speeds. Hence, flares of the type discussed above may be accompanied by additional radiation processes caused by this energetic plasma. Recent multiwavelength observations of Mrk 501 argue that high energy processes in blazars are complicated (see e.g. Pian et al.~1998, Protheroe et al.~1998), and more than one different mechanism may be responsible for producing the observed photon spectrum. The motion of plasma, and additionally the longitudinal component of the magnetic field in the jet (we do not include its influence on the final $\\gamma$-ray spectrum) could provide collimation of the produced radiation in the direction perpendicular to the disk surface. In conclusion, this model predicts that two types of high energy flares may occur if the particles (hadrons, leptons) are accelerated to high energies in reconnection regions on the accretion disk surface. Pure $\\gamma$-ray flares can be caused either by secondary $e^\\pm$ pairs produced during acceleration of protons which collide with the soft disk photons, or by primary leptons accelerated in the reconnection region. In this latter case, the $\\gamma$-ray spectrum can extend up to TeV energies if the disk temperature is low enough. For large, luminous disks, the magnetic field energy can be efficiently transferred to relativistic protons which can initiate $\\gamma$-ray and neutrino flares via production of pions in collisions with soft disk photons. Therefore, we suggest that blazars of FSRQ type, for which $\\gamma$-ray spectra show some evidence of a cut-offs at higher energies (Pohl et al.~1997) and features of large accretion disks (emission lines), may become sources of GeV $\\gamma$-rays and $\\sim 10^{17}$ eV neutrinos. In contrast, BL Lac objects, with spectra extending to TeV energies, should not emit high energy neutrinos according to the present model." + }, + "9802/astro-ph9802077_arXiv.txt": { + "abstract": "This review concentrates on nucleosynthesis processes in general and their applications to massive stars and supernovae. A brief initial introduction is given to the physics in astrophysical plasmas which governs composition changes. We present the basic equations for thermonuclear reaction rates and nuclear reaction networks. The required nuclear physics input for reaction rates is discussed, i.e. cross sections for nuclear reactions, photodisintegrations, electron and positron captures, neutrino captures, inelastic neutrino scattering, and beta-decay half-lives. We examine especially the present state of uncertainties in predicting thermonuclear reaction rates, while the status of experiments is discussed by others in this volume (see M. Wiescher). It follows a brief review of hydrostatic burning stages in stellar evolution before discussing the fate of massive stars, i.e. the nucleosynthesis in type II supernova explosions (SNe II). Except for SNe Ia, which are explained by exploding white dwarfs in binary stellar systems (which will not be discussed here), all other supernova types seem to be linked to the gravitational collapse of massive stars (M$>$8M$_\\odot$) at the end of their hydrostatic evolution. SN1987A, the first type II supernova for which the progenitor star was known, is used as an example for nucleosynthesis calculations. Finally, we discuss the production of heavy elements in the r-process up to Th and U and its possible connection to supernovae. ", + "introduction": "In this section we want to outline the essential features of thermonuclear reaction rates and nuclear reaction networks. This serves the purpose to define a unified terminology to be used throughout the review, more detailed discussions can be found in Fowler, Caughlan, \\& Zimmerman (1967,1975), Clayton (1983), Rolfs \\& Rodney (1988), Thielemann, Nomoto, \\& Hashimoto (1994), and Arnett (1996). \\subsection{Thermonuclear Reaction Rates} The nuclear cross section for a reaction between target $j$ and projectile $k$ is defined by \\begin{equation} \\sigma = {\\rm{number\\ of\\ reactions\\ target^{-1} sec^{-1}} \\over {flux\\ of\\ incoming\\ projectiles}} = {{r/n_j} \\over {n_k v}}. \\end{equation} The second equality holds for the case that the relative velocity between targets with the number density $n_j$ and projectiles with number density $n_k$ is constant and has the value $v$. Then $r$, the number of reactions per cm$^3$ and sec, can be expressed as $r=\\sigma v n_j n_k$. More generally, when targets and projectiles follow specific distributions, $r$ is given by \\begin{equation} r_{j,k}=\\int \\sigma \\vert \\vec v_j -\\vec v_k\\vert d^3 n_j d^3 n_k. \\end{equation} The evaluation of this integral depends on the type of particles and distributions which are involved. For nuclei $j$ and $k$ in an astrophysical plasma, obeying a Maxwell-Boltzmann distribution, \\begin{equation} d^3n_j=n_j ({{m_j} \\over {2\\pi kT}})^{3/2} {\\rm exp}(- {{m_jv_j^2} \\over {2kT}}) d^3v_j, \\end{equation} Eq.(1.2) simplifies to $r_{j,k}=<\\sigma v> n_j n_k$. The thermonuclear reaction rates have the form (Fowler, Caughlan, \\& Zimmerman 1967, Clayton 1983) \\begin{subeqnarray} r_{j,k} & = &<\\sigma v>_{j,k} n_j n_k \\\\ :& = &<\\sigma v>_{j,k}=({8 \\over {\\mu \\pi}})^{1/2} (kT)^{-3/2} \\int_0 ^\\infty E \\sigma (E) {\\rm exp}(-E/kT) dE. \\end{subeqnarray} Here $\\mu$ denotes the reduced mass of the target-projectile system. In astrophysical plasmas with high densities and/or low temperatures, effects of electron screening become highly important. This means that the reacting nuclei, due to the background of electrons and nuclei, feel a different Coulomb repulsion than in the case of bare nuclei. Under most conditions (with non-vanishing temperatures) the generalized reaction rate integral can be separated into the traditional expression without screening [Eq.(1.4)] and a screening factor (see e.g. Salpeter \\& van Horn 1969, Itoh, Totsuji, \\& Ichimaru 1977, Hansen, Torrie, \\& Veillefosse 1977, Alastuey \\& Jancovici 1978, Itoh et al. 1979, Ichimaru, Tanaka, Iyetomi 1984, Ichimaru \\& Utsumi 1983, 1984, Thielemann \\& Truran 1987, Fushiki \\& Lamb 1987, Itoh et al. 1990, Schramm \\& Koonin 1990, Ichimaru 1993, Chabrier \\& Schatzman 1994, Kitamura \\& Ichimaru 1995, Brown \\& Sawyer 1997) \\begin{equation} ^*=f_{scr}(Z_j,Z_k,\\rho,T,Y_i) . \\end{equation} This screening factor is dependent on the charge of the involved particles, the density, temperature, and the composition of the plasma. Here $Y_i$ denotes the abundance of nucleus $i$ defined by $Y_i=n_i/(\\rho N_A)$, where $n_i$ is the number density of nuclei per unit volume and $N_A$ Avogadro's number. At high densities and low temperatures screening factors can enhance reactions by many orders of magnitude and lead to {\\it pycnonuclear ignition}. In the extreme case of very low temperatures, where reactions are only possible via ground state oscillations of the nuclei in a Coulomb lattice, Eq.(1.5) breaks down, because it was derived under the assumption of a Boltzmann distribution (for recent references see Fushiki \\& Lamb 1987, Itoh et al. 1990, Schramm \\& Koonin 1990, Ichimaru 1993, Chabrier \\& Schatzman 1994, Ichimaru 1996). When in Eq.(1.2) particle $k$ is a photon, the relative velocity is always c and quantities in the integral are not dependent on $d^3n_j$. Thus it simplifies to $r_j=\\lambda_{j,\\gamma} n_j$ and $\\lambda_{j,\\gamma}$ results from an integration of the photodisintegration cross section over a Planck distribution for photons of temperature $T$ \\begin{subeqnarray} d^3n_\\gamma & = & {{1} \\over {\\pi^2 (c\\hbar)^3}} {{E_\\gamma^2}\\over { {\\rm exp}(E_\\gamma /kT)-1}} dE_\\gamma\\\\ r_j & = & \\lambda_{j,\\gamma} (T)n_j= {{\\int d^3n_j}\\over {\\pi ^2 (c \\hbar )^3}} \\int_0 ^\\infty {{c \\sigma(E_\\gamma ) E_{\\gamma}^2} \\over {{\\rm exp}(E_\\gamma /kT) -1}} dE_\\gamma. \\end{subeqnarray} There is, however, no direct need to evaluate photodisintegration cross sections, because, due to detailed balance, they can be expressed by the capture cross sections for the inverse reaction $l+m\\rightarrow j+\\gamma$ (Fowler et al. 1967) \\begin{equation} \\lambda_{j,\\gamma} (T)= ({{G_l G_m} \\over G_j}) ({{A_l A_m} \\over A_j})^{3/2} ({{m_ukT} \\over {2\\pi \\hbar^2}})^{3/2} {\\rm exp} (-Q_{lm}/kT). \\end{equation} This expression depends on the reaction Q-value $Q_{lm}$, the temperature $T$, the inverse reaction rate $$, the partition functions $G(T)=\\sum_i (2J_i+1){\\rm exp}(-E_i/kT)$ and the mass numbers $A$ of the participating nuclei in a thermal bath of temperature $T$. A procedure similar to Eq.(1.6) is used for electron captures by nuclei. Because the electron is about 2000 times less massive than a nucleon, the velocity of the nucleus $j$ is negligible in the center of mass system in comparison to the electron velocity ($\\vert \\vec v_j- \\vec v_e \\vert \\approx \\vert \\vec v_e \\vert$). The electron capture cross section has to be integrated over a Boltzmann, partially degenerate, or Fermi distribution of electrons, dependent on the astrophysical conditions. The electron capture rates are a function of $T$ and $n_e=Y_e \\rho N_A$, the electron number density (Fuller, Fowler, \\& Newman 1980, 1982, 1985). In a neutral, completely ionized plasma, the electron abundance is equal to the total proton abundance in nuclei $Y_e=\\sum_i Z_i Y_i$ and \\begin{equation} r_j=\\lambda_{j,e} (T,\\rho Y_e)n_j. \\end{equation} The same authors generalized this treatment for the capture of positrons, which are in a thermal equilibrium with photons, electrons, and nuclei. At high densities ($\\rho >10^{12}$gcm$^{-3}$) the size of the neutrino scattering cross section on nuclei and electrons ensures that enough scattering events occur to thermalize a neutrino distribution. Then also the inverse process to electron capture (neutrino capture) can occur and the neutrino capture rate can be expresses similar to Eqs.(1.6) or (1.8), integrating over the neutrino distribution (e.g. Fuller \\& Meyer 1995). Also inelastic neutrino scattering on nuclei can be expressed in this form. The latter can cause particle emission, like in photodisintegrations (e.g. Woosley et al. 1990, Kolbe et al. 1992, 1993, 1995, Qian et al. 1996). It is also possible that a thermal equilibrium among neutrinos was established at a different location than at the point where the reaction occurs. In such a case the neutrino distribution can be characterized by a chemical potential and a temperature which is not necessarily equal to the local temperature. Finally, for normal decays, like beta or alpha decays with half-life $\\tau_{1/2}$, we obtain an equation similar to Eqs.(1.6) or (1.8) with a decay constant $\\lambda_j=\\ln 2/\\tau_{1/2}$ and \\begin{equation} r_j=\\lambda_j n_j. \\end{equation} \\subsection{Nuclear Reaction Networks} The time derivative of the number densities of each of the species in an astrophysical plasma (at constant density) is governed by the different expressions for $r$, the number of reactions per cm$^3$ and sec, as discussed above for the different reaction mechanisms which can change nuclear abundances \\begin{equation} ({{\\partial n_i} \\over {\\partial t}})_{\\rho =const}= \\sum_j N^i _j r_j + \\sum_{j,k} N^i _{j,k} r_{j,k} + \\sum_{j,k,l} N^i _{j,k,l} r_{j,k,l}. \\end{equation} The reactions listed on the right hand side of the equation belong to the three categories of reactions: (1) decays, photodisintegrations, electron and positron captures and neutrino induced reactions ($r_j=\\lambda_j n_j$), (2) two-particle reactions ($r_{j,k}=n_j n_k$), and (3) three-particle reactions ($r_{j,k,l}= n_j n_k n_l$) like the triple-alpha process, which can be interpreted as successive captures with an intermediate unstable target (see e.g. Nomoto, Thielemann, \\& Miyaji 1985, G\\\"orres, Wiescher, \\& Thielemann 1995). The individual $N^{i}$'s are given by: $N^i_j = N_i$, $N^i_{j,k} = N_i / \\prod_{m=1}^{n_m} | N_{j_m} |! $, and $N^i_{j,k,l} = N_i / \\prod_{m=1}^{n_m} |N_{j_m}|!$. The $N_i's$ can be positive or negative numbers and specify how many particles of species $i$ are created or destroyed in a reaction. The denominators, including factorials, run over the $n_m$ different species destroyed in the reaction and avoid double counting of the number of reactions when identical particles react with each other (for example in the $^{12}$C+$^{12}$C or the triple-alpha reaction; for details see Fowler et al. 1967). In order to exclude changes in the number densities $\\dot n_i$, which are only due to expansion or contraction of the gas, the nuclear abundances $Y_i =n_i/(\\rho N_A)$ were introduced. For a nucleus with atomic weight $A_i$, $A_iY_i$ represents the mass fraction of this nucleus, therefore $\\sum A_iY_i=1$. In terms of nuclear abundances $Y_i$, a reaction network is described by the following set of differential equations \\begin{equation} \\dot Y_i = \\sum_j N^i _j \\lambda_j Y_j + \\sum_{j,k} N^i _{j,k} \\rho N_A Y_j Y_k + \\sum_{j,k,l} N^i _{j,k,l} \\rho^2 N_A^2 Y_j Y_k Y_l. \\end{equation} Eq.(1.11) derives directly from Eq.(1.10) when the definition for the $Y_i's$ is introduced. This set of differential equations is solved with a fully implicit treatment. Then the stiff set of differential equations can be rewritten (see e.g. Press et al. 1986, \\S 15.6) as difference equations of the form $\\Delta Y_i/\\Delta t=f_i(Y_j(t+\\Delta t))$, where $Y_i(t+\\Delta t)=Y_i(t)+\\Delta Y_i$. In this treatment, all quantities on the right hand side are evaluated at time $t+\\Delta t$. This results in a set of non-linear equations for the new abundances $Y_i(t+\\Delta t)$, which can be solved using a multi-dimensional Newton-Raphson iteration procedure. The total energy generation per gram, due to nuclear reactions in a time step $\\Delta t$ which changed the abundances by $\\Delta Y_i$, is expressed in terms of the mass excess $M_{ex,i}c^2$ of the participating nuclei (Audi \\& Wapstra 1995) \\begin{subeqnarray} \\Delta \\epsilon & = & - \\sum_i \\Delta Y_i N_A M_{ex,i}c^2 \\\\ \\dot \\epsilon & = & - \\sum_i \\dot Y_i N_A M_{ex,i}c^2. \\end{subeqnarray} As noted above, the important ingredients to nucleosynthesis calculations are decay half-lives, electron and positron capture rates, photodisintegrations, neutrino induced reaction rates, and strong interaction cross sections. Beta-decay half-lives for unstable nuclei have been predicted by Takahashi, Yamada, \\& Kondo (1973), Klapdor, Metzinger, \\& Oda (1984), Takahashi \\& Yokoi (1987, also including temperature effects) and more recently with improved quasi particle RPA calculations (Staudt et al.~1989, 1990, M\\\"oller \\& Randrup 1990, Hirsch et al.~1992, Pfeiffer \\& Kratz 1996, M\\\"oller, Nix, \\& Kratz 1997, Borzov 1996, 1997). Electron and positron capture calculations have been performed by Fuller, Fowler, \\& Newman (1980, 1982, 1985) for a large variety of nuclei with mass numbers between A=20 and A=60. For revisions see also Takahara et al. (1989) and for heavier nuclei Aufderheide et al.~(1994), Sutaria, Sheikh, \\& Ray (1997). Rates for inelastic neutrino scattering have been presented by Woosley et al. (1990) and Kolbe et al. (1992, 1993, 1995). Photodisintegration rates can be calculated via detailed balance from the reverse capture rates. Experimental nuclear rates for light nuclei have been discussed in detail in the reviews by Rolfs, Trautvetter, \\& Rodney (1987), Filippone (1987), the book by Rolfs \\& Rodney (1988), the recent review on 40 years after B$^2$FH by Wallerstein et al. (1997), and the NuPECC report on nuclear and particle astrophysics (Baraffe et al. 1997). The most recent experimental charged particle rate compilations are the ones by Caughlan \\& Fowler (1988) and Arnould et al. (1997). Experimental neutron capture cross sections are summarized by Bao \\& K\\\"appeler (1987, 1997), Beer, Voss, \\& Winters (1992), and Wisshak et al.~(1997). Rates for unstable (light) nuclei are given by Malaney \\& Fowler (1988, 1989), Wiescher et al.~(1986, 1987, 1988ab, 1989ab, 1990), Thomas et al.~(1993,1994), van Wormer et al. (1994), Rauscher et al. (1994), and Schatz et al. (1997). For additional information see the article by M. Wiescher (this volume). For the vast number of medium and heavy nuclei which exhibit a high density of excited states at capture energies, Hauser-Feshbach (statistical model) calculations are applicable. The most recent compilations were provided by Holmes et al. (1975), Woosley et al. (1978), and Thielemann, Arnould, \\& Truran (1987, for a detailed discussion of the methods involved and neutron capture cross sections for heavy unstable nuclei see also section 3.4 and the appendix in Cowan, Thielemann, Truran 1991). Improvements in level densities (Rauscher, Thielemann, \\& Kratz 1997), alpha potentials, and the consistent treatment of isospin mixing will lead to the next generation of theoretical rate predictions (Rauscher et al. 1998). Some of it will be discussed in the following section. ", + "conclusions": "" + }, + "9802/astro-ph9802063_arXiv.txt": { + "abstract": "Gravitational lensing can be used to analyze the redshift distribution of faint galaxies. In particular the magnification bias modifies locally the galaxy number density of lensed sources observed in lensing clusters. This {\\sl depletion area} probes the redshift distribution of galaxies beyond $B=25$. In this proceedings I present this new tool to infer the redhsift distribution of faint galaxies. ", + "introduction": "With the coming of 10 meter class telescopes equipped with wide field multi-object spectrographs, deep redshift surveys will be extended to thousands of galaxies and will permit to explore in detail the evolution of clustering of galaxies, the history of star formation up to $z=4$ for galaxies with $B \\le 25$.\\\\ The study of galaxies with magnitudes $B>25$ are also important for the models of galaxy formation : we do not know yet whether they are all at large redshift or if there is a significant fraction of faint nearby dwarfs galaxies. The knowledge of their redshift distribution is also necessary for mass reconstruction using lensing inversion, and can be a major source of uncertainty in the mass determination for the most distant lensing clusters (see Luppino \\& Kaiser 1997). Bernardeau et al (1997) and Jain \\& Seljak (1997) have emphasized that even the study the large-scale mass distribution using weak lensing need the redshift distribution of the faintest galaxies, because the variance and the skewness of the magnification strongly depends of the redshift of the lensed background sources. Unfortunately, beyond $B=25$, even 10 meter class telescopes are unable to provide redshifts of a complete sample of galaxies. The possibility of using photometric redshifts which was proposed by the beginning of eighties is now re-investigated in great details. But observations as well as reliability tests are still underway (Connolly et al. 1995. Since they are based on theoretical evolution scenarios of galaxies, their predictions about faintest galaxies are not fully confirmed yet. Furthermore, there is no hope to calibrate the photometric redshifts of the faint samples with spectroscopic data. An attractive alternative to spectroscopy consists in using of the magnification and distortion effects induced by gravitational lensing on extended objects. In particular, the magnification bias can eventually produces depletion areas in the projected galaxy number density of background sources observed in rich clusters whose size and shape depend on their redshift distribution. In the following section I present the basic principle of the technique and first results. ", + "conclusions": "The redshift distribution of galaxies beyond $B=25$ is a crucial scientific question for galaxy evolution and weak lensing studies for mass reconstruction. The depletion curves of galaxy number density produced by magnification bias is an innovative way which can probe the redshift distribution of galaxies as faint as $B=28$. The first tentative by Fort et al (1997) demonstrates that depletion curves can be observed in Cl0024+1654 and A370. However, a good modeling of the lensing clusters is needed in order to infer the redshift distribution of the lensed sources. This method is still at its infancy and the first results are questionable. Hence, it must be considered jointly with other techniques like photometric redshifts or lensing inversion (Kneib et al 1994, 1996). \\\\ Whatever the method, how can we be sure that these redshifts obtained from non-standard and indirect techniques are correct ? Preliminary deep spectroscopic and multicolor photometric surveys of arclets show that the faintest galaxies seem to have a redshift distribution like the ones predicted by Fort et al (Pell\\'o, private communication). But this key issue demands ultra-deep CCD spectroscopic exposures with the VLTs. This should be in the future a major challenge for the gravitational telescopes." + }, + "9802/astro-ph9802255_arXiv.txt": { + "abstract": "We investigate Big Bang nucleosynthesis (hereafter, BBN) in a cosmic environment characterised by a distribution of small-scale matter--antimatter domains. Production of antimatter domains in a baryo-asymmetric universe is predicted in some electroweak baryogenesis scenarios. We find that cosmic antimatter domains of size exceeding the neutron-diffusion length at temperature $T \\sim 1$~MeV significantly affect the light-element production. Annihilation of antimatter preferentially occurs on neutrons such that antimatter domains may yield a reduction of the $^4$He abundance relative to a standard BBN scenario. In the limiting case, all neutrons will be removed before the onset of light-element production, and a universe with net baryon number but without production of light elements results. In general, antimatter domains spoil agreement between BBN abundance yields and observationally inferred primordial abundances limits which allows us to derive limits on their presence in the early universe. However, if only small amounts of antimatter are present, BBN with low deuterium {\\it and} low $^4$He, as seemingly favored by current observational data, is possible. ", + "introduction": " ", + "conclusions": "" + }, + "9802/astro-ph9802125_arXiv.txt": { + "abstract": "We point out that already existing literature on relativistic collisionless MHD shocks show that the parameter $\\sigma\\equiv$ upstream proper magnetic energy density/upstream rest mass energy density, plays an important role in determining the structure and accelerating properties of such shocks. By adopting the value of $\\sigma\\approx 0.002$ which corresponds to the relativistic shock associated with the Crab nebula, and by using appropriate relativistic shock jump conditions, we obtain here a generous upper-limit on the value of (proper) magnetic field, $B_{\\rm sh}\\approx 1.5 \\times 10^{-3} \\eta n_1^{1/2} $G, for gamma ray burst (GRB) blast wave. Here, $\\eta \\equiv E/Mc^2$, where $E$ is the energy and $M$ is the mass of the baryons entrained in the original fireball (FB), and $n_1$ is the proper number density of the ambient medium. Further, we point out that, in realistic cases, the actual value $B_{\\rm sh}$ could be as low as $\\sim 5\\times 10^{-6} \\eta n_1^{1/2}$G. realistic cases. ", + "introduction": "Understanding the phenomenon of GRBs is one of the important problems of recent astrophysics. Fortunately, following the discovery of a cosmological redshift in the May 08, 1997 event, it is certain now that some or all of them could be of cosmological origin Metzger et al. (1997). Whether cosmological or galactic, GRB phenomenon is broadly understood in terms of a standard model developed by Cavallo \\& Rees (1978), Goodman (1986), Paczynski (1986), Eichler et al. (1989), Shemi \\& Piran (1990). Nonetheless, as far as the origin of the complex nonthermal observed GRB spectra are concerned, an important development took place with the work of Rees \\& Meszaros (1992) and Meszaros \\& Rees (1993) suggesting that the cosmic fireballs (FB) with an optimal amount of baryonic pollution, $\\eta \\sim 10^2 -10^3$, could explain such spectra, where $\\eta \\equiv E/Mc^2$, where $E$ is the energy and $M$ is the mass of the baryons entrained in the original fireball (FB). Meszaros \\& Rees suggested that as the baryon polluted FB deposits half of its original momentum onto the ambient medium, presumably, the bare interstellar medium (ISM), at $r=r_{\\rm d}$, the so-called deceleration radius, the blast wave becomes very strong and radiates part of its energy. For further appreciation of this paper it would be appropriate to crudely visualize the geometry associated with the blast wave in terms of a 1-D simple diagram (Mitra 1998, henceforth M98). Here region (1) is the ambient ISM, the lab frame, $S_1$ is the forward shock moving ahead of the contact discontinuity $S$, the location of the original FB boundary. Region (2) represents the (forward) shocked fluid and it is this region which is the site for the particle acceleration and gamma ray production in this standard model. The region (4) is the unperurbed FB and (3) is the part of the FB compressed by the reverse shock front $S_2$. It was shown in M98 that, in the context of this standard model, the reverse shock plays an insignificant role in the overall energy balance and may be neglected for dynamical purposes. The gamma rays are likely to be produced either by a synchrotron process or a self-synchrotron-Compton process occuring near the region (2) and the most crucial factor for the success of such processes is the value of the comoving magnetic field $B_2'=B_{\\rm sh}$. Here prime denotes respective comoving quantities, i.e. respective {\\em proper values}, whereas `*' denotes quantities measured in the rest frame of the forward shock $S_1$. The question of probable generation of a magnetic field in a relativistic (or nonrelativistic) shock is a poorly understood topic, and, practically, most of the authors have therefore been compelled to use an equiparition argument to estimate the same (Meszaros \\& Rees 1993, Cheng \\& Wei 1996, Vietri 1995, Waxman 1995): \\begin{equation} {B_{\\rm sh}^2\\over 8 \\pi} \\sim \\eta^2 n_1 m c^2 \\end{equation} But by recalling the basic definition of \\begin{equation} r_{\\rm d} \\approx \\left({3 E \\over 2\\pi c^2 \\eta^2 m n_1}\\right)^{1/3} \\approx 7\\times 10^{15} E_{51}^{1/3} \\eta_3^{-2/3} n_1^{-2/3}~{\\rm cm} \\end{equation} where $\\eta_3= 10^{-3} \\eta$ and $E_{51}= E/10^{51}$erg s$^{-1}$, it can be easily verified that the energy density shown on the R.H.S. of eq.(1) directly corresponds to the region (4), i.e, the unperturbed FB, having a proper density (M98) \\begin{equation} n_4'= {E(\\gamma_{\\rm F}/ \\eta) \\over 4\\pi r^3 c^2 m} \\approx 5\\times 10^7 E_{51} r_{15}^{-3} (\\gamma_{\\rm F}/\\eta)~ {\\rm cm}^{-3}, \\end{equation} i.e., actually, \\begin{equation} {B_4'^2\\over 8 \\pi} \\approx \\gamma_{\\rm F}^2 n_1 m c^2 \\end{equation} Here $\\gamma_{\\rm F}$ is the bulk Lorentz factor (LF) of $S_1$ in the lab frame (1), $m$ is the mass of a proton, and $n_1$ is the particle number density of the ambient medium in units of 1$cm^{-3}$. Note that for the region (1) the comoving frame coincides with the lab frame and at $r=r_{\\rm d}$, we will have $\\gamma_{\\rm F} \\approx \\eta/2$ in the Meszaros \\& Rees scenario. On the other hand, we need to apply the equipartition argument in the downstream of the shock, i.e., in the region, which is expected to be turbulent, and which is, in any case, the site for the particle acceleration. And, it will be seen in the next section that, equipartition argument yields approximately the same value of $B_2'$ in the shocked fluid, and this is important, because, energy density, in itself, is not a Lorentz invariant quantity. Nevertheless, the question we want to pose here is how justified is this assumption of equipartition in the context of GRBs and whether by adopting this brute assumption we are running into conflict with some well established feature of relativistic collisionless shocks. As was stressed in Mitra (1996), equipartition, as a general physical concept may be found to be valid in steady-state situations like the ISM where the plasma interacts with the particles and currents over astronomically significant time scales. As to dynamic situations, there are hints that many young supernova remnants are endowded with freshly generated magnetic fields which are considerably higher than the bare ISM values $\\sim 3\\times 10^{-6}$G. Nonetheless, even in such cases, the age of the supernova could be thousands of years and the enhanced magnetic field is usually much smaller than what is obtained by naive equipartition arguments. In fact, it was clearly anticipated by Meszaros, Rees, \\& Papathanssiou (1994) that the equipartition argument can at best serve as a broad guide to determine the actual value of $B_{\\rm sh}$, and accordingly, they introduced a {\\em completely free parameter}, $\\lambda \\le 1$, tagged onto the naively obtained value of $B_{\\rm sh}^{\\rm eq}$: \\begin{equation} B_{\\rm sh} \\sim 4\\times 10^2 n_1^{1/2} \\eta_3 \\lambda^{1/2} ~{\\rm G} \\end{equation} Further, it could be possible to apply the basic equipartition idea at $t=0$ to the initial FB, and then evaluate the value of the instantaneous $B_{\\rm FB}$ or $B_{\\rm sh}$ by using the flux-freezing condition. And again, in this case, if we symbolize our ignorance through the free parameter, $\\xi \\le 1$, it follows that (Meszaros, Rees \\& Papathanassiou 1994) \\begin{equation} B_{\\rm sh} \\sim 0.4 \\xi^{1/2} E_{51}^{-1/6} n_1^{2/3} \\eta_3^2 ~{\\rm G} \\end{equation} Thus, in this paper, we would attempt to invoke a known feature of relativistic collisionless MHD shocks to obtain physically significant upper limits on $B_{\\rm sh}$. ", + "conclusions": "Having obtained this generous upper-limit let us now ponder how justified we are in adopting a value of $\\sigma$ appropriate for Crab. Remember that the Crab shock is practically a standing shock and the upstream region is not the bare ISM. On the other hand, the {\\em upstream comprises the plasma ejected by the Crab pulsar during its life time of nearly thousand years}. It is very much likely that, in the past, the value of $\\sigma$ for Crab was much lower, and the present value has been slowly built up over these thousand years. In contrast the case of the GRB blast wave ploughing through the bare ISM is quite different, and, it is highly improbable that, within a lab frame time scale of few seconds or less, the upstream medium ahead of the shock front (as seen by in the lab frame) can raise its magnetic field from a value of $\\sim 3 \\times 10^{-6}$G to $\\sim 9\\times 10^{-3}$G. We feel that, instead, the following scenario is more plausible: The value of $B_1'$ probably remains close to its unperturbed value $\\sim 3\\times 10^{-6}$G; however the shock could be near perpendicular resulting in a value of $B_{\\rm sh} \\sim 4 \\gamma_{12} B_1'\\sim \\sqrt{2} \\eta B_1' \\sim 4 10^{-3} \\eta_3$G. It is also probable, there may not be any stable shock formation at all invalidating the rigid relations between $B_1'$ and $B_2'$ employed so far, and on the other hand there may be instantaneous spikes in the downstream magnetic field (Langdon, Arons, \\& Max 1988) : \\begin{equation} {B_{2\\rm max}^* \\over B_{\\rm sh}^*} \\approx \\left[ 1+ (2/\\sigma) \\right]^{1/2} \\end{equation} and which may erratically raise the shocked field to a $B_{\\rm max} \\sim 0.1 -1$G. However, formation of a shock-like discontinuity, either steady or fluctuating requires that the ambient medium should be such that the {\\em leading particles} of the FB, i.e., the piston driving the shock, either individually or collectively impart significant amount of their momentum on the ambient medium. We endeavoured to examine this critical but usually overlooked problem in M96. Noting the similarity between the present problem and the one involving propagation of high energy cosmic rays in the ISM, and also recalling that phenomenologically and observationally obtained parameter, the {\\em spatial diffusion coefficient} describes the entire collective interaction of the cosmic rays and the ISM, we found, that, it is implausible that the FB can produce any thing akin to a GRB blast wave in the ISM (M96). This is so for the simple reason that the BF-ISM interaction time scale estimated in this way could be as large as $\\sim 10^7$s ! Probably, we may also examine here whether, the electrons or the left over pairs of the FB, rather than the protons, might carry out the job of imparting energy onto the ambient medium (B. Paczynski, private communication). For a given saturation bulk LF of the FB ($\\gamma_{\\rm F} \\approx \\eta/2$), the electrons or positrons will be less energetic by a factor of $\\sim 2000$. And, if we are assuming the Bohm limit of the diffusion coefficient, the deflection length as well as the lab frame deflection time will accordingly be smaller by a factor of $\\sim 2000$. If the numer of leptons in the FB could overwhelm the protons by a factor of $\\sim 2000$, there would have been equal amount of energy residing in protons and leptons, and, the foregoing reduction in time scale would have really meant that the FB-ISM interaction time would be reduced by the same factor. Unfortunately, this is not the case, the number of leptons per proton at the saturation stage of the optimally baryon polluted FB is indeed $\\gtrsim 1$. This means that the eventual FB-ISM interaction time scale should be what was obtained in M96, i.e., \\begin{equation} t_i \\sim 3 \\times 10^4 \\eta_3 ~{\\rm s} \\end{equation} unless it is found that the lepton-ion energy transfer time scale is smaller than this above time scale. It may be probable that at best there may be soliton like discontinuites. At any rate, we do not think, we have the final answer of such questions at this moment, and we realize that, we are simply attempting to understand various aspects subject to our (present author's) available knowledge. The upshot of this discussion is that the original GRB is likely to be produced either by internal collisions within the FB as has long been suspected (Paczynski \\& Xu, Rees \\& Meszaros 1994, Papathanssiou \\& Meszaros 1996) or if the original FB is propagating within a medium which has a modest baryonic mass but is dense enough to absorb the FB momentum either by binary collisions or by collisionless MHD process. Finally, reverting back to the original mandate of this paper, in the case of a supposed GRB blast wave propagating in the ISM and even ignoring the disturbing possibility that no shock like discontinuity may be formed on the GRB time scale, we feel that, the maximum value of the magnetic field in the shocked fluid may not exceed $\\sim 1 $G. We would like to emphasize here the fact that this conclusion does not at all imply that the FB can not excite a blast wave (initially relativistic and then non-relativistic) on a much larger time scale of days or weeks and which is necessary for explaining the GRB after glow in various low energy brackets. \\newpage" + }, + "9802/astro-ph9802313_arXiv.txt": { + "abstract": "The study of the very young open cluster NGC 6231 clearly shows the presence of a mass segregation for the most massive stars. These observations, combined with those concerning other young objects and very recent numerical simulations, strongly support the hypothesis of an initial origin for the mass segregation of the most massive stars. These results led to the conclusion that massive stars form near the center of clusters. They are strong constraints for scenarii of star and stellar cluster formation. ", + "introduction": "The aim of the present work is to give observational constraints to better understand the dynamical evolution of open clusters (OCs). For an analysis of this evolution the observation and study of clusters of different ages is required. In order to fix the \"initial conditions\" we chose to extensively observe NGC 6231, which is a rich and very young OC: it has an age of 3-4 Myr and contains more than one hundred O and B stars (Raboud 1996, Raboud et al. 1997). The data collected allowed us to complete the analysis of this cluster structure (Raboud 1997, Raboud \\& Mermilliod 1998, RM98). ", + "conclusions": "The above results allow us to propose a qualitative scenario for the evolution of mass segregation with age in OCs (RM98): \\bigskip \\textbf{(I)} The most massive stars form in the center of clusters. Several hypotheses could be made to explain this phenomenon: dynamical friction between protostellar clouds and inter-protostellar medium (Larson 1991, Gorti \\& Bhatt 1995, 1996); collision and coalescence of protostellar clouds (Murray \\& Lin 1996); the accretion of matter during stellar formation phases. This accretion could be faster in regions of higher temperature and turbulence (Maeder 1997), i.e. in the center of protocluster clouds, thus leading to the formation of more massive stars in these regions. This last hypothesis implies that the IMF is dependent on the local physical conditions. In the context of massive star formation in the center of clusters, it is worth noting that we observe numerous examples of multiple systems of O-stars in the center of very young OCs. In the case of NGC 6231, 8 stars among the 10 brightest are spectroscopic binaries with periods shorter than 6 days. Moreover, we observe trapezium systems of O-stars in the ONC, NGC 6823 and Tr 37. Four-component and triple systems have also been found in NGC 2362 and Collinder 228. \\textbf{(II)} In less than $10^{7}$ yr these spatially concentrated massive stars will disappear due to stellar evolution. As they represent a non-negligible percentage of the total mass of the cluster (between $\\sim$10 and 30 \\% in the case of NGC 6231), the disappearance of these massive stars could lead to a violent relaxation phase. If a mass segregation was previously established in the cluster it could be more or less erased during this phase, depending on the importance of the initial population of massive stars. We are then {\\it possibly} left with a cluster presenting {\\it no} mass segregation at all. NGC 6531 (Forbes 1996) provides an example of such a cluster: it is 8 $\\times$ 10$^{6}$ yr old and does not contain any stars with masses greater than 20 M$_\\odot$, which make up the concentrated population in NGC 6231. Forbes shows convincingly that NGC 6531 does not exhibit any mass segregation, and he explains his observation by the young age of the cluster. According to him, NGC 6531 is too young for dynamical evolution to have left any significant impression. But this hypothesis was based on an estimation of $t_{r}$ and suffers the drawbacks described in the Discussion. Another interesting point related to the disappearance of the massive stars is the stability of the cluster. It is possible that a bound cluster becomes unbound after this violent phase. Numerical simulations by Terlevich (1987) show that clusters with flat initial mass functions have to be rich enough to survive the initial violent period of mass loss. \\textbf{(III)} The last point of our scenario is that all mass segregation observed in older clusters is \\textit{merely} the consequence of the cluster's \\textit{dynamical evolution}. \\bigskip To better quantify this hypothesis of a possible double origin (initial and dynamical) of the mass segregation we need to analyse the structure of OCs old enough (around 10$^{7}$ yr) to have lost their most massive stars. Thus, one consequence of our hypothesis is that some of these clusters, those which initially contained an important population of massive stars, should not present any mass segregation." + }, + "9802/astro-ph9802196_arXiv.txt": { + "abstract": "Emission of a charged particle propagating in a medium with a curved magnetic field is reconsidered stressing the analogy between this emission mechanism and collective Cherenkov-type plasma emission. It is explained how this mechanism differs from conventional Cherenkov, cyclotron or curvature emission and how it includes, to some extent, the features of each of these mechanisms. Presence of a medium supporting subluminous waves is essential for the possibility of wave amplification by particles streaming along the curved magnetic field with a finite curvature drift. We suggest an analogy between the curvature drift emission and the anomalous cyclotron-Cherenkov emission. Treating the emission in cylindrical coordinates in the plane-wave-like approximation allows one to compute the single particle emissivity and growth rate of the Cherenkov-drift instability. We compare the growth rates calculated using the single particle emissivity and using the dielectric tensor of one dimensional plasma streaming along the curved field. In calculating the single particle emissivity it is essential to know the normal modes of the medium and their polarization which can be found from the dielectric tensor of the medium. This emission mechanism may be important for the problem of pulsar radio emission generation. ", + "introduction": "Studies of relativistic strongly magnetized plasma in astrophysical setting (like pulsar magnetosphere) have shown the possibility of a new mechanism of electromagnetic wave generation. This new mechanism, which combines features of conventional Cherenkov, cyclotron and curvature radiation, deserves more detailed consideration from the fundamental physics point of view. Besides investigating the physical nature of this process, we also reconcile different available approaches to this problem. In this work we discuss this novel emission mechanism of a charged particle streaming with relativistic velocity along curved magnetic field line in a medium. A weak inhomogeneity of the magnetic field results in a curvature drift motion of the particle perpendicular to the local plane of the magnetic field line. A gradient drift (proportional to ${\\bf \\nabla \\cdot B}$) is much smaller than the curvature drift and will be neglected. When the motion of the particle parallel to the magnetic field is ultrarelativistic the drift motion even in the weakly inhomogeneous field can become weakly relativistic resulting in a new type of generation of {\\it electromagnetic}, vacuumlike waves. Presence of three ingredients ( strong but finite magnetic field, inhomogeneity of the field and a medium with the index of refraction larger than unity) is essential for the emission. We will call this mechanism Cherenkov-drift emission stressing the fact that microphysically it is virtually Cherenkov-type emission process. Conventional consideration of the curvature emission (\\cite{Blandford1975}, \\cite{ZheleznyakovShaposhnikov}, \\cite{MelroseLou}, \\cite{Melrosebook1}) emphasize the analogy between curvature emission and conventional cyclotron emission. To our opinion this approach, though formally correct, has limited applicability and misses some important physical properties of the emission mechanism. In a separate approach Kazbegi {\\it et al.} \\cite{Kazbegi} considered this process calculating a dielectric tensor of the inhomogeneous magnetized medium, thus treating the emission process as a collective effect. They showed that maser action is possible only if a medium supports subluminous waves. In this work we show how these two approaches can be reconciled and argue that the dielectric tensor approach, which treats the Cherenkov-drift emission as a collective process, has a wider applicability. The interplay between cyclotron (or synchrotron) and Cherenkov radiation has been a long-standing matter of interest. Schwinger {\\it et al.} \\cite{Schwinger} discussed the relation between these two seemingly different emission mechanisms. They showed that conventional synchrotron emission and Cherenkov radiation may be regarded as respectively limiting cases of $|n-1|\\, \\ll \\,1 $ and $B=\\,0$ of a synergetic (using the terminology of Schwinger {\\it et al.} \\cite{Schwinger}) cyclotron-Cherenkov radiation. In another work \\cite{LyutikovCher} this analogy has been discussed to include cyclotron-Cherenkov emission at the anomalous Doppler resonance. An important new aspect of our work (as compared with \\cite{Schwinger} and \\cite{LyutikovCher} ) is that we take into account inhomogeneity of the medium. Physical origin of the emission in the case of Cherenkov-type and synchrotron-type processes is quite different. In the case of Cherenkov-type process the emission may be attributed to the electromagnetic polarization shock front that develops in a dielectric medium due to the passage of a charged particle with speed larger than phase speed of waves in a medium. It is virtually a collective emission process. In the case of synchrotron-type process, the emission may be attributed to the Lorentz force acting on a particle in a magnetic field. Cherenkov-type emission is impossible in vacuum and in a medium with the refractive index smaller than unity. Cyclotron emission at the anomalous Doppler effect (cyclotron-Cherenkov emission) is an interesting example of Cherenkov-type emission process of a particle in a magnetic field. It is impossible in vacuum and requires a superluminous motion of a particle along the magnetic field. Thus, the emission at the anomalous Doppler effect is attributed to the polarization shock front that a spiraling particle induces in a medium. In our opinion, the Curvature-drift emission may be viewed as a Cherenkov-type emission that bears the same relation to the conventional curvature emission as the cyclotron emission at the anomalous Doppler effect bear to the conventional cyclotron emission. In this work we consider a Curvature-drift emission of the particles in the ground gyrational state. It is possible to obtain the emissivity for particles in excited gyrational state by the method of dielectric tensor \\cite{Akhalkatsi1997}. Then, conventional cyclotron (cyclotron-Cherenkov), Cherenkov and curvature emission mechanisms may be viewed as corresponding limits of the Cherenkov-drift mechanism in the cases of homogeneous magnetic field (in a medium), medium without magnetic field, and inhomogeneous magnetic field without a medium. In Section \\ref{Model1} we discuss our set up of the problem and how it differs from the previous consideration. In Section \\ref{emissivity} we calculate a single particle emissivity of the Cherenkov-drift mechanism and find the growth for kinetic beam instability toward excitation of electromagnetic waves at the Cherenkov-drift resonance. Then the results are compared with those obtained by the dielectric tensor method. ", + "conclusions": "In this paper we considered a new Cherenkov-drift emission mechanism that combines features of the conventional cyclotron, Cherenkov and curvature emission. We argued, that from the microphysical point of view this emission mechanism may be regarded as a Cherenkov-type process in inhomogeneous magnetic field. Considering emission process in cylindrical coordinates we have obtained the single particle emissivities. We also pointed out, that in order to obtain correct expressions for the emissivities it is necessary to use the polarization vectors of the normal modes of the medium. Finally, we calculated the growth rates of the Cherenkov-drift instability in a strongly magnetized electron-positron plasma." + }, + "9802/astro-ph9802019_arXiv.txt": { + "abstract": "We present an analysis of compact star clusters in deep HST/WFPC2 images of NGC 1275. B and R band photometry of roughly 3000 clusters shows a bimodality in the B-R colors, suggesting that distinct old and young cluster populations are present. The small spread in the colors of the blue clusters is consistent with the hypothesis that they are a single age population, with an inferred age of 0.1 to 1 Gyr. The luminosity function shows increasing numbers of blue clusters to the limit of our photometry, which reaches several magnitudes past the turnover predicted if the cluster population were identical to current Galactic globulars seen at a younger age. The blue clusters have a spatial distribution which is more centrally peaked than that of the red clusters. The individual clusters are slightly resolved, with core radii $\\ltaprx 0.75$ pc if they have modified Hubble profiles. We estimate the specific frequencies of the old and young populations and discuss the uncertainties in these estimates. We find that the specific frequency of the young population in NGC 1275 is currently larger than that of the old population and will remain so as the young population evolves, even if the majority of the low mass clusters are eventually destroyed. If the young population formed during a previous merger, this suggests that mergers can increase the specific frequency of globulars in a galaxy. However, the presently observed young population likely contains too few clusters to have a significant impact on the overall specific frequency as it will be observed in the future. ", + "introduction": "Early HST observations of NGC 1275, the central galaxy in the Perseus cluster, revealed a population of about 60 blue (V-R $\\sim$ 0.3) star clusters surrounding the nucleus (Holtzman \\etal \\markcite{hol92} 1992). The color of these clusters suggests an age of roughly 300 million years based on the models of Charlot and Bruzual \\markcite{cb91} (1991). Spectra of the brightest object (Zepf \\etal \\markcite{zep95} 1995) also suggest an age of 0.1 to 0.9 Gyr, based on a comparison of the observed line widths to those predicted by the models of Bruzual and Charlot \\markcite{bc93} (1993). None of the clusters seem to have H$\\alpha$ emission, with the exception of one object found by Shields \\& Filippenko \\markcite{sf90} (1990), which appears to be a much younger object. The blue clusters in the original WFPC1 images appear unresolved, suggesting sizes of less than 15 parsecs. The brightest object has V = 18.9, which corresponds to M$_V$ = -15.8 for $H_0$ = 75 km/s/Mpc and cz=5264 km/s (Strauss \\etal \\markcite{str92} 1992). The brightnesses of the blue clusters suggest masses of between 2 $\\times$ 10$^4$ and 1 $\\times$ 10$^8$ M$_{\\odot}$, depending on the assumed age and distance, if one assumes that the objects are star clusters with a Salpeter IMF. The observed sizes, luminosities, and the estimated masses suggest that these objects may be young analogues of globular clusters. In the past several years, massive young clusters have been observed in a variety of other galaxies. Lutz \\markcite{lut91} (1991) detected young globular cluster candidates in a ground-based study of the merger remnant NGC 3597, and these have been confirmed to be compact by HST observations (Holtzman \\etal \\markcite{hol96} 1996). Candidate young globular clusters have also been found in other interacting systems, including NGC 7252 (Whitmore \\etal \\markcite{whit93} 1993), NGC 4038/9 (Whitmore \\& Schweizer \\markcite{ws95} 1995), NGC 3921 (Schweizer \\etal \\markcite{sch96} 1996), among others. A few massive clusters are seen in the starburst galaxy NGC 253 (Watson \\etal \\markcite{wat96} 1996), while others have been detected in the ring galaxies NGC 1097 and NGC 6951 (Barth \\etal \\markcite{bhfs95} 1995) and the dwarf galaxies NGC 1569, NGC 1705 (O'Connell \\etal \\markcite{ogh94} 1994), and He 2-10 (Conti \\& Vacca \\markcite{cv94} 1994). While no generally accepted picture has emerged as to what conditions lead to the formation of young clusters, galaxy interactions appear to be an important component. It is particularly difficult to determine the precise mechanism responsible for the presence of these objects in NGC 1275 because of the myriad peculiarities of this galaxy, including the presence of a significant amount of dust, streamers of H$\\alpha$ emission, an active nucleus, and the location of the galaxy at the center of a cooling flow in the Perseus cluster. Two hypotheses have been offered for the origin of these clusters, namely, that they were formed from the substantial mass deposition of the Perseus cluster cooling flow (200 M$_{\\odot}$/yr) or that the cluster formation was triggered by a galaxy-galaxy interaction. Holtzman \\etal \\markcite{hol92} (1992) preferred the latter hypothesis based on the observed lack of spread in the WFPC1 colors, which implies a common age for the objects, and the appearance of a ripple in the galaxy light which suggests that a previous interaction may have occurred. Richer \\etal \\markcite{ric93} (1993) found larger color spreads in high resolution CFHT images and preferred the cooling flow hypothesis. However, Holtzman \\etal \\markcite{hol96} (1996) did not detect young clusters in three of a sample of four other cooling flow galaxies; the central cluster galaxy in Abell 1795 may have a few clusters, but it also has a peculiar morphology which suggests a previous interaction. The merger hypothesis is particularly interesting in light of the theory that elliptical galaxies form through mergers and the observation that ellipticals have higher specific frequencies of globular clusters than spirals, where the specific frequency is a measure of the number of globular clusters per unit luminosity of the host galaxy. Ashman \\& Zepf \\markcite{az92} (1992) investigate the proposition that clusters might form during mergers and suggest that, in such an event, cluster formation would occur in a brief burst, resulting in a set of newly formed clusters with a common age and color. This would lead to a cluster system with a bimodal distribution of cluster colors reflecting the difference in metallicity or age (or both) between the original and the newly formed clusters. They predict that the spatial distribution of the younger clusters would be more sharply peaked toward the center of the galaxy than that of the old cluster system, because the old globular cluster populations of the two progenitor galaxies would remain spatially extended and probably be dynamically heated during the merger, while the new clusters would be formed out of gas which becomes more centrally concentrated during a merger. However, as Harris \\etal \\markcite{har95} (1995) note, an increase in specific frequency as a result of a merger requires not only that globular clusters form during such an event, but that they form {\\it preferentially} over non-cluster stars as compared to the ratio of clusters to background stars in the progenitor galaxies. Few estimates of the specific frequency of these young cluster systems exist. Watson \\etal \\markcite{wat96} (1996) suggest that the 4 young clusters in NGC 253 most likely have formed with a large specific frequency. In the merger remnant NGC 3921 Schweizer \\etal \\markcite{sch96} (1996) find that the blue clusters will increase the overall number of clusters enough that the galaxy will come to have the specific frequency of an elliptical within 7 Gyr. Miller \\etal \\markcite{mil97} (1997) find that the specific frequency of globular clusters in the merger remnant NGC 7252 will rise over the next 15 Gyr as the background population fades to resemble that of an elliptical. These results all suggest that interactions can increase the specific frequency in a galaxy. Luminosity functions of old globular cluster systems have been well studied and used as a part of the extragalactic distance scale because of their uniformity from galaxy to galaxy; they are roughly Gaussian in shape with a peak near M$_V$ = -7.3 (Harris \\markcite{har96} 1996). The luminosity functions of most of the recently discovered young cluster systems are poorly determined because of either small number statistics (few young clusters in the galaxy) or incompleteness. A notable exception is the young cluster system in NGC 4038/9, which shows no turnover to two magnitudes fainter than the turnover predicted for a typical old cluster system, even after allowing for the expected fading of the clusters based on stellar population models. Based on this, van den Bergh \\markcite{vdb95} (1995) argues that this cluster system may be intrinsically different from old globular cluster systems. However, Mateo \\markcite{ma93} (1993) notes that, at least in the LMC, the observed increasing cluster luminosity function is well modelled by a combination of globular and open clusters. Also, several authors have recently suggested that substantial numbers of clusters may be destroyed over a Hubble time (Gnedin \\& Ostriker \\markcite{go97} 1997; Elmegreen \\& Efremov \\markcite{ee97} 1997); mechanisms include evaporation and tidal disruption by galactic bulges and disks. We have obtained WFPC2 images of NGC 1275 which go about 4 magnitudes deeper (in the red) than the WFPC1 observations of Holtzman et al. \\markcite{hol92} (1992). These observations provide more accurate colors than previous observations and probe the cluster population to approximately 2.5 magnitudes fainter than the turnover expected if this population is identical to the Galactic globular cluster system seen at an younger age. Section 2 briefly summarizes the observations and the reductions. Section 3 discusses our analysis and potential sources of error. In section 4 we present our results including the photometry, luminosity function of the clusters, surface density distribution, an estimate of the specific frequency, and analysis of the sizes of the objects. ", + "conclusions": "From new deep observations of NGC 1275 with HST, we identify roughly 3000 objects which appear to be compact star clusters. The color distribution of the cluster system is bimodal, with a blue population which has (B-R)$_0$ $\\sim$ 0.4 and a red population which has (B-R)$_0$ $\\sim$ 1.3. We suggest that the red objects are members of the old globular cluster system and that the blue objects are members of a young globular cluster system. In an apparently dust-free region, the spread in the colors of the blue clusters is small enough that it is entirely attributable to scatter from errors in the photometry, mostly due to errors in sky subtraction. This suggests that the blue clusters are a single color, single age population. This argues against the hypothesis that clusters have been forming continuously from the cooling flow, and supports the hypothesis that their formation may have been triggered by a previous merger. The luminosity function continues to rise to the limit of our observations and is inconsistent with a Gaussian globular cluster luminosity function peaking near M$_V$ = -7.3 after correction for evolutionary effects. A similar luminosity function observed in NGC 4038/9 (the Antennae) has been used by van den Bergh \\markcite{vdb95} (1995) to argue that the young clusters observed there are not true globular clusters. However, the masses and sizes of the individual young clusters appear to be comparable to those of globulars. We suggest that either the luminosity function evolves, with fainter clusters being preferentially destroyed as time passes, or that the initial luminosity/mass function of the young cluster system is different from that of typical old globular cluster systems. If cluster destruction is responsible for the difference in luminosity function, a destruction mechanism which preferentially destroys fainter, lower mass clusters must be invoked. In order to create a turnover in the luminosity function at the same mass as in old cluster systems, the destruction of over 90\\% of the clusters currently seen in NGC 1275 would be required, given our assumption of 4 magnitudes of fading. Given these likely differences in the luminosity functions of different globular cluster systems, whether caused by evolutionary effects or by initial differences, it is clear that some caution must be exercised in using the cluster luminosity function as a distance indicator. We have attempted to estimate the specific frequency of the young population seen in NGC 1275 to determine whether it is likely that the overall specific frequency as seen at a future time could be increased because of the proposed merger-related formation of globulars. We note that the overall specific frequency of NGC 1275 will only be slightly increased by the proposed merger event which formed the blue clusters, due to the small number of clusters expected to survive for a Hubble time. But we estimate a large specific frequency of the merger related population, which suggests that mergers are efficient in forming globular clusters. However, corrections for incompleteness at faint magnitudes, clusters residing outside the current field of view, uncertainties in the mass and luminosity of the stellar population formed during a merger event, and, especially, corrections for evolutionary effects in both the background and cluster populations are all significant effects for which we have only relatively crude estimates." + }, + "9802/astro-ph9802109_arXiv.txt": { + "abstract": "Observations of large scale structure (LSS) and the Cosmic Microwave Background (CMB) each place separate constraints on the values of cosmological parameters. We calculate a joint likelihood based on various CMB experiments and the IRAS 1.2Jy galaxy redshift survey and use this to find an overall optimum with respect to the free parameters. Our formulation self-consistently takes account of the underlying mass distribution, which affects both the CMB potential fluctuations and the IRAS redshift distortion. This not only allows more accurate parameter estimation, but also removes the parameter degeneracy which handicaps calculations based on either approach alone. The family of Cold Dark Matter (CDM) models analysed corresponds to a spatially-flat universe with an initially scale-invariant spectrum and a cosmological constant. Free parameters in the joint model are the mass density due to all matter ($\\omegam$), Hubble's parameter ($h = H_0 / 100\\kms {\\rm Mpc}^{-1}$), the quadrupole normalisation of the CMB power spectrum ($Q$) in $\\mu$K, and the IRAS light-to-mass bias ($\\bi$). Throughout the analysis, the baryonic density ($\\Omega_b$) is required is to satisfy the nucleosynthesis constraint $\\Omega_b h^2 = 0.024$. Results from the two data sets show good agreement, and the joint optimum lies at $\\omegam = 0.39$, $h = 0.53$, $Q = 16.96$ $\\mu$K, and $\\bi = 1.21$. The 68 per cent confidence intervals are: $0.29<\\omegam<0.53$, $0.393\\times 10^{10}\\lsun$ (Fig.~1). The observed X-ray spectra of galaxies with high $\\lx/\\lb$ ratios are consistent with thermal emission from hot, optically thin gas, while those of low $\\lx/\\lb$ objects can be mostly accounted for by emission from stellar sources (Kim, Fabbiano, \\& Trinchieri 1992). The scatter in the $\\lx -\\lb$ diagram has been recognized as the most striking feature of the X-ray properties of early-type galaxies. Using new apparent magnitudes and fundamental plane distances, it was shown that this scatter is reduced by 20\\%, but not eliminated (Donnelly, Faber \\& O'Connell 1990; see Fig.~1b). This result is based on the old estimate of the X-ray fluxes (that of Canizares, Fabbiano, \\& Trinchieri 1987), and on just half of the final sample of X-ray galaxies produced by Fabbiano et al. 1992 (see Eskridge, Fabbiano \\& Kim 1995 for a detailed comparison of the statistical results by Donnelly et al. with those obtained using the whole sample). One can argue that the large dispersion in $\\lx $ is definitively not the result of distance errors on the basis of the fact that a scatter of the same size as in Fig.~1a is present even in the distance-independent diagram of $\\lx/\\lb$ versus the central stellar velocity dispersion (e.g., Eskridge et al. 1995). Many theoretical models were developed to explain the findings above, including numerical simulations of the behavior of gas flows fed by stellar mass loss and heated by type Ia supernovae (SNIa). Steady state cooling flow models were investigated first (Nulsen et al. 1984, \\cite{sw87},1988), and it was found that these can only reproduce X-ray bright galaxies. Evolutionary models with a SNIa rate approximately constant with time have been carried out by Mathews \\& Loewenstein (1986), Loewenstein \\& Mathews (1987), and David, Forman, \\& Jones (1990, 1991). After a very brief initial wind phase, driven by the explosion of type II supernovae, the resulting flow evolution goes from a global inflow to a wind, which is experienced only by the smallest galaxies by the present time. David \\etal (1991) conclude that all galaxies above $\\lb\\simeq 3\\times 10^{10}\\lsun$ host a cooling flow. So, as in the steady state cooling flow scenario, the scatter in $\\lx $ at fixed $\\lb$ has to be explained by a combination of environmental differences (\\cite{ws91}) and by large variations from galaxy to galaxy in the stellar mass loss rate per unit $\\lb$, the efficiency of thermalization of the stellar mass loss, and that of thermal instabilities in the hot gas. An alternative way of explaining the scatter in the $\\lx -\\lb$ diagram is given by the evolutionary scenario of D'Ercole \\etal (1989), and Ciotti \\etal (1991, CDPR). Assuming that the SNIa explosion rate is declining with time slightly faster than the rate with which mass is lost by the stars, in the beginning the energy released by SNIa's can drive the gas out of the galaxies through a supersonic wind. As the collective SNIa energy input decreases, a subsonic outflow takes place, which gradually slows until a central cooling catastrophe leads to the onset of an inflow. At fixed $\\lb$, any of the three phases wind, outflow or inflow can be found at the present epoch, depending only on the various depths and shapes of the potential well of the galaxies. In this way both the large scatter in $\\lx $ and the trend in the spectral properties are accounted for at the same time: in the X-ray bright galaxies the soft X-ray emitting gas dominates the emission, being in the inflow phase, that resembles a cooling flow; in the X-ray faint galaxies the hard stellar emission dominates, these being in the wind phase. Recent $ASCA$ observations seems to indicate that the SNIa's activity is suppressed in early-type galaxies, which implies that the CDPR scenario is essentially ruled out. Since this issue is far from closed (\\S 1.3), it is still worthwile to explore the effects of SNIa's on the flows, using the updated rate given by recent optical surveys; this is lower than that adopted by CDPR, which was 0.67 the standard one estimated by Tammann (1982). This aspect, together with the need for changes in other galaxy properties crucial for the evolution of hot gas flows, are discussed below. \\begin{figure}[htbp] \\picplace{10cm} \\caption{The $\\lx-\\lb$ diagram of early-type galaxies with X-ray emission detected by the {\\it Einstein} satellite; X-ray fluxes are from the final catalog of Fabbiano \\etal (1992). Apparent B magnitudes and distances are from Fabbiano et al. (1992) in Fig.~1a, and from Donnelly et al. (1990) in Fig.~1b (see Section 1). The dashed line gives an estimate for the stellar source contribution to $\\lx$ (from Kim \\etal 1992). Also shown with triangles are the positions of the models calculated here (see Section 3.5), to which the stellar source contribution has been added.} \\label{Fig1} \\end{figure} \\subsection {New stellar density profiles} Previous numerical simulations of hot gas flows used the King (1972) stellar density distribution mainly for computational ease. This distribution has an inner region of constant density (the so called ``core'') of the order of a few hundred parsecs, which keeps the time step of the numerical simulation reasonably small (a stellar density increase implies a reduction of the characteristic hydrodynamical time step). Another advantage of density profiles with a constant density core is that the central regions can be resolved using a larger grid size; this again allows a larger time step (see CDPR for a more quantitative discussion). It is well known though that a much better description of the surface brightness profiles of ellipticals is given by the de Vaucouleurs (1948) law. A very good fit to this law is given by the Jaffe (1983) and by the Hernquist (1990) distributions, that have the advantage that all their dynamical properties can be expressed analytically. These distributions belong to the family of the so called $\\gamma$-models, that has been widely explored recently (\\cite{d93}; \\cite{tremetal94}): \\begin{equation} \\rho (r)=M{(3-\\gamma)\\over 4\\pi} {\\rc\\over r^{\\gamma}(\\rc+r)^{4-\\gamma}}. \\end{equation} The density profile of the Hernquist law has $\\gamma=1$, while that of the Jaffe law has $\\gamma=2$. Inside $\\rc$ the density of $\\gamma$-models increases as $r^{-\\gamma}$, a significant difference with respect to King models: over a few hundreds of parsecs at the center the $\\gamma$-models are power laws. The existence of cores of constant surface brightness has been definitively ruled out by ground based observations (Lauer 1985, Kormendy 1985), and recently by the Hubble Space Telescope (Jaffe et al. 1994; \\cite{lauetal95}), that has shown how the central surface brightness profile is described by a power law as far in as can be observed, i.e., $\\sim 10$ pc in Virgo. From the point of view of a very accurate photometry, the surface brightness law implied by eq. (1) cannot reproduce well both the envelope and the very center of all the ellipticals observed by HST. The Jaffe law, though, gives a general description of ellipticals accurate enough for the treatment of hot gas flows, on the scales that are relevant for the problem (from a few tens to several thousands of parsecs). \\subsection {New dark matter estimates and distributions} Attempts to estimate the amount of nonluminous mass in elliptical galaxies have been made recently through extensive searches for dynamical evidence of dark matter, either with systematic observations of ionized gas probing the gravitational field (\\cite{pizetal97}) or with measurements of stellar velocity dispersions profiles out to large radii (e.g., \\cite{berteal94}; Carollo et al. 1995). These optical studies are confined to within one or two effective radii $\\rref$, and typically find that luminous matter dominates the mass distribution inside $\\sim\\rref$, while dark matter begins to be dynamically important at 2--3 $\\rref$. In particular, for a sample of X-ray emitting galaxies, it has been found that dark matter halos are not much more massive than the luminous component, with the dark-to-luminous mass ratio $\\mh/\\ms\\approx 1-6$; the value of $\\mh/\\ms\\sim 2$ is most common (\\cite{saglieal93}). X-ray emission from hot gas provides a great potential for mapping the mass of ellipticals to larger distances (e.g., Fabian et al. 1986). The standard method employed derives from the equation of hydrostatic equilibrium, and requires the knowledge of the gas temperature profile. Attempts to apply this technique to the {\\it Einstein} data yielded a much larger component of dark matter than found from optical data, but these results are very uncertain because temperature profiles are poorly determined (Forman, Jones, \\& Tucker 1985; \\cite{fab89}). Using superior X-ray data provided by {\\it ROSAT} and {\\it ASCA}, Buote \\& Canizares (1997) proved that mass does not follow the optical light, out to many $\\rref$, in NGC720 and NGC1332. Adopting plausible gas and mass models, they find $\\mh/\\ms >3$ for NGC1332, and $\\mh/\\ms >7$ in NGC720 at 90\\% confidence; $\\mh $ prevails exterior to $\\rref$. Similarly Mushotzky \\etal (1994) derive $\\mh/\\ms \\sim 8$ within $8\\rref$ for NGC4636, and analogous results have been obtained from {\\it ROSAT} data of NGC507 and NGC499 by Kim \\& Fabbiano (1995). {\\it AXAF} will have the combined spatial and spectral resolution to measure accurately the presence of different spectral components, their relative flux and their spatial distribution, and to produce more accurate mass distribution from X-rays. The radial density distribution of the dark haloes of ellipticals is not well constrained by observations; theoretical arguments favor a peaked profile (Ciotti \\& Pellegrini 1992; Evans \\& Collett 1997), and high resolution numerical simulations of dissipationless collapse produce a density distribution with $\\gamma\\simeq 1$ near the center (\\cite{dc91}; Navarro, Frenk, \\& White 1996; \\cite{w96}; Fukushige \\& Makino 1997, and references therein). Previous works studying the hot gas flow evolution always used quasi isothermal haloes at least nine times more massive than the luminous component. We are motivated now to explore even the effects produced by dark haloes not as massive as supposed before, and not quasi-isothermal. \\subsection{New Type Ia supernova rates} Nearby SNIa rates in early-type galaxies have been carefully reanalyzed recently, and this important ingredient of the simulations of hot gas flows has been revised. From optical surveys it was estimated to be 0.88$\\,h^2$ SNu (\\cite{tam82}), and then 0.98$\\,h^2$ SNu (\\cite{vdbt91}). Most recent estimates agree on lower values: 0.25$\\,h^2$ SNu (\\cite{vdbmc94}) and 0.24$\\,h^2$ SNu (Turatto, Cappellaro, \\& Benetti 1994), where $h=H_{\\circ}/100$ and 1 SNu = 1 SNIa per century per $10^{10}L_{\\rm B\\odot}$. In principle, constraints on the SNIa rate can be given also by estimates of the iron abundance in galactic flows (see, e.g., \\cite{rcdp93}, and references therein). These were first attempted using data from the {\\it Ginga} satellite for NGC4472, NGC4636, NGC1399 (Ohashi \\etal 1990; Awaki \\etal 1991; Ikebe \\etal 1992), and then from the {\\it BBXRT} satellite for NGC1399, NGC4472 (\\cite{serleal93}), and more recently from {\\it ASCA} with a superior spectral energy resolution (\\cite{loeweal94}; Awaki \\etal 1994; Arimoto \\etal 1997; Matsumoto et al. 1997). Under the assumption of solar abundance ratio, the analysis of all these data suggests a very low iron abundance, consistent with no SNIa's enrichment and even lower than that of the stellar component. However, some authors have found that more complex multi-temperature models with higher abundance give a better fit of the data (\\cite{kimeal96}, Buote \\& Fabian 1997). Moreover, the above results are based on iron line diagnostic tools whose reliability has been questioned (e.g., \\cite{arieal97}), especially because of uncertainties affecting the Fe L-shell atomic physics for temperatures less than 2 keV, that are typical of hot gas flows in ellipticals (Liedahl, Osterheld, \\& Goldstein 1995). Line emission even in simple (e.g., isothermal) astrophysical plasmas needs to be understood, and reliable fits to the data made, before secure consequences concerning the SNIa rate can be drawn from X-ray determined abundances. \\medskip In summary, recent optical studies agree on a present epoch SNIa rate much lower than previously used, and indicate that also the dark matter content could be lower. Moreover, cuspier density profiles, especially for the stellar distribution, should be used. All this raises the question of whether the CDPR scenario is altered by these changes in the main ingredients of the problem, and more in general what are the effects on the properties of hot gas flows. In this paper we address these points, with a new set of hydrodynamical simulations. In Section 2 we present the galaxy models, the source terms, and the integration techniques. In Section 3 we discuss the main properties of the evolution of gas flows in our new models, and we compare them with the observations and the CDPR results. In Section 4 we discuss the results using energetic arguments, and in Section 5 the main conclusions are summarized. ", + "conclusions": "Using evolutionary hydrodynamical simulations, we have investigated the properties of hot gas flows in elliptical galaxies described by a Jaffe stellar density profile, and containing various amounts of dark matter, distributed as a Hernquist law; the rate of SNIa explosions is one-fourth of the Tammann's rate, or lower. The results can be summarized as follows. \\smallskip The simulations show that most of the gas flows are strongly decoupled, i.e., they develop an inflow in the central region of the galaxy, while the external parts are still degassing\\footnote{Note that decoupled flows are present also in flat and rotating models with low $\\vt$, albeit for reasons different from those causing the PWs studied here (D'Ercole \\& Ciotti 1997).}. The stagnation radius of these flows (PWs) may range from a small fraction of the optical effective radius to many $\\rref$. This situation is present even when the global energy balance of the flow indicates that the energy inputs from SNIa's and from the thermalization of the stellar random motions is high enough to unbind all the gas ($\\chi <1$). Alternatively, the gas can be outflowing from the outer part of the galaxy even when the global energy balance indicates that the energy available is less than needed to unbind it all ($\\chi >1$). So, $\\chi$ is not a good indicator of the flow phase; larger $\\chi $ values, though, correspond to larger stagnation radii, for a fixed $\\lb$. Global inflows are produced when the dark matter content is so high to bind the gas over the whole galaxy. The X-ray luminosity of PWs is higher when $\\vt$ is lower, because the inflow region is larger; for global inflows $\\lx$ is higher when $\\vt$ is also higher, because the gas is hotter. The key factor causing large $\\lx$ variations in JH models is the size of the central inflow region. The lowest $\\lx/\\lb$ values observed can be easily reproduced by PWs; high dark matter contents are required to approach the highest $\\lx$ observed. The shape of the X-ray surface brightness profile of PWs is close to that of global inflows, over most of the optical image, if the stagnation radius is a few $\\rref$; it is externally less peaked than that, if $\\rst$ is a fraction of $\\rref$, i.e., when the gas is hotter than in the inflow case. \\smallskip The strong decoupling of the flow is explained by the radial trend of the local energy balance $\\chil$, that is very peaked for steep mass density profiles as in the JH models. Previously used density profiles which were flat at the center -- such as the King models plus quasi-isothermal dark halos -- showed a larger tendency to have global flow phases, because also their $\\chil (r)$ is flatter. So, the change of the mass profile has an important effect: both numerical simulations and analytical calculations show how peaked density profiles preferentially develop decoupled gas flows (in the sense that the region of the parameter space populated by PWs is large compared to that of inflows or winds). \\smallskip The results of our simulations compare with the observations in two main aspects. The first is that most of the observed spread in $\\lx$ shown by the data is easily reproduced: the highest $\\lx$ are again, as in CDPR, associated with global inflows, while the bulk of the galaxies are now in PW, rather than in outflow; global winds are present only in the smallest galaxies. The second is that the presence of cold gas at the center of Es, and a peaked X-ray surface brightness profile, cannot be unequivocally associated to a cooling flow: a PW could be present as well, with a significant part of the galaxy degassing. Particularly, a galaxy can have $\\vt=0.25$ and still a non negligible amount of cold matter at the center. In addition, since $\\rst$ varies over a wide range of radii, PWs can accumulate largely varying quantities of cold gas." + }, + "9802/astro-ph9802203_arXiv.txt": { + "abstract": " ", + "introduction": "The solar system is believed to reside in a very hot (temperature T $ \\sim ~ 10^6 $ K) and tenuous (electron density $ n_e ~ \\approx ~ 0.005 ~ cm^{-3} $) X-ray emitting cavity, which is typical of the coronal phase of the ISM, possibly produced by supernovae or winds from hot, massive stars (Cox \\& Reynolds 1987; McCammon \\& Sanders 1990). This region, often termed as the Local Bubble, has been of considerable interest to both observers and theorists alike owing to its proximity and atypical nature. The X-ray data reveal that the Local Bubble has an elongated geometry, extending up to 200$-$300 pc perpendicular to the galactic plane and up to 50$-$100 pc in the plane (Snowden et al. 1990). This region was also noted for its deficiency of neutral hydrogen gas (Paresce 1984; Frisch \\& York 1983), and recent studies in the extreme-ultraviolet (EUV) waveband of a large sample of bright sources (Warwick et al. 1993) estimate the gas density to be $ \\sim $ 0.05 $ {\\rm cm ^{-3}} $. The local interstellar medium (LISM), which consists of the bubble and its surroundings, has rather loose definitions (Cox \\& Reynolds 1987; Bochkarev 1987) and in this paper we use this term to mean a region of a few 100 pc surrounding us. Though adequate understanding exists to say that the LISM deviates significantly in its properties from the galactic average, there are several aspects that lack satisfactory understanding, specially its detailed morphological characteristics, accurate size, nature of the boundary and properties of the material in it. Propagation effects on radio waves emitted by pulsars, such as dispersion and scattering, probe the distribution of thermal plasma in the interstellar medium (ISM). In particular, observable effects of scattering of radio waves from pulsars, enable us to probe the electron density fluctuations in the ISM (see Rickett (1990) for a review). These density fluctuations are thought to arise from turbulence and hence scintillation studies of pulsars also provide information on the nature and distribution of the interstellar turbulence (Cordes, Weisberg \\& Boriakoff 1985). The distribution of electron density fluctuations in the LISM can differ significantly from the typical ISM due to the Local Bubble. It is reasonable to expect the Local Bubble and its environment to play a substantial role in determining the scintillation properties of nearby (distance \\la 1 kpc) pulsars. Such pulsars therefore form potential tools to study the LISM. Pulsar scintillation properties are best studied using their dynamic spectra $-$ records of intensity variations in the time-frequency plane $-$ which shows random intensity modulations that fade over narrow frequency ranges and short time intervals. The phenomenon giving rise to this effect, known as diffractive scintillation, arises from scattering of pulsar signals by the small-scale ($ 10^7 ~ \\la $ s $ \\la ~ 10^9 $ m) density fluctuations present in the ISM (Rickett 1990; Cordes, Pidwerbetsky \\& Lovelace 1986; Spangler 1988). The decorrelation bandwidth, {\\it i.e.}, the frequency range over which the intensity decorrelates, provides information on the strength of scattering along the line-of-sight. In addition, the average scintillation properties of pulsar dynamic spectra are also expected to show variations over time scales of days to weeks at metre wavelengths as a result of refractive scintillation effects (Cordes et al. 1986; Romani, Narayan \\& Blandford 1986; Rickett 1990). These arise from propagation of pulsar signals through large-scale (s $ \\ga ~ 10 ^{11} $ m) density irregularities in the ISM. Therefore, estimates of average scintillation properties from a few epochs of observations of dynamic spectra are prone to errors due to refractive scintillation effects. A large number of long stretches of data taken over time spans of months to years are essential to get reliable and accurate estimates of scintillation parameters. Systematic studies of this kind have not been carried out for many pulsars and most earlier measurements were based on a few epochs of observations. Not much is known about the electron density distribution of the interstellar gas in the LISM. Current models predict uniform electron density and uniform electron density fluctuations, except for a smooth dependence with height above the galactic plane (Taylor \\& Cordes 1993, Cordes et al. 1985). Recent scintillation observations of PSR B0950+08 suggest that the interior of the bubble is dominated by relatively lower magnitudes of density fluctuations (Phillips \\& Clegg 1992). However, there has been no accurate and systematic study of the connection between scintillation properties of local pulsars and the structure of the LISM. During 1993$-$95 we carried out a long-term, systematic study of the scintillation properties of twenty nearby pulsars, with the two-fold aim of (i) studying refractive scintillation effects and (ii) understanding the LISM. In this paper, we describe the results of the second study. The results of the first part are presented in another paper (Bhat, Rao \\& Gupta 1997c, under preparation). Our observations are described in Section 2. In Section 3, we present the results of the data analysis and discuss its implications on the distribution of electron density fluctuations in the LISM. In Section 4, we discuss the viability of simple models of specific density structures in explaining the present observations. The success and limitations of our best model are briefed in Section 5, where we also compare it with pertinent results from other published studies on the LISM. At the end, we suggest some possible tests for the present model and some useful observations which will improve upon the present understanding of the LISM. ", + "conclusions": "For the first time, the structure of the LISM has been modelled using the results from a systematic, long-term pulsar scintillation study. Our analysis based on the scintillation properties of twenty nearby pulsars suggests that the large-scale distribution of the ionized material in the solar neighborhood is not uniform. Systematic trends have been seen in the scattering properties of pulsars, which imply a coherent structure of electron density fluctuations in the LISM. The detailed analysis of the observed anomalous scattering effects shows that such a structure is highly asymmetric relative to the location of the Sun. Simple models in which the solar neighborhood has an enhanced or a reduced scattering strength relative to the ambient medium fail to reproduce the scattering anomalies. To explain our observations, we need a three-component scattering medium in which the solar neighborhood is surrounded by a shell of much higher density fluctuations embedded in the normal, large scale ISM. We are also able to put reasonable constraints on the geometrical and the scattering properties characterizing the size, location and density fluctuations of such a structure. The shell has an ellipsoidal morphology and is much more extended away from the galactic plane than in the plane, with radii of $ \\sim $ 270$-$330 pc and $ \\sim $ 60$-$75 for the sections along the galactic poles and through the galactic plane respectively. The Sun is located away from the centre by about $ \\sim $ 20$-$35 pc. The density fluctuations in the shell are much larger than those in the interior and in the outer region and there is a suggestion that it decreases with the height above the galactic plane. We also find that the morphology of our scattering structure is similar to that of the Local Bubble known from various earlier studies based on HI, X-ray and UV data. The LISM and the distribution of electron density fluctuations in it are likely to be more complex than that is suggested by our simplified model, but we hope the present work will serve as a useful framework within which more detailed questions can be addressed. { {\\it Acknowledgments:} The authors would like to thank J. Chengalur and M. Vivekanand for reading the manuscript and giving useful comments. We also thank M. Vivekanand for providing the software for pulsar data acquisition, and V. Balasubramanian for the telescope time and technical help with the observations. We thank the referee, S. R. Spangler, for several fruitful comments and suggestions, which improved the contents of the paper. }" + }, + "9802/astro-ph9802173_arXiv.txt": { + "abstract": "We have used the Optical Redshift Survey (ORS; \\cite{san95}) to construct the gravity field due to fluctuations in the galaxy density field out to distances of $8000\\kms$. At large scales where linear theory applies, the comparison of this gravity field with the observed peculiar velocity field offers a powerful cosmological probe, because the predicted flow field is proportional to the parameter $\\Omega^{0.6}/b$, where $\\Omega$ is the matter density and $b$ is the bias of the galaxy distribution. The more densely sampled ORS gravity field, to excellent approximation, matches that of the earlier \\iras\\ 1.2-Jy redshift survey (\\cite{fis95a}), provided $\\beta$ is reduced by a factor $b_{opt}/b_{IRAS} \\approx 1.4.$ Apart from this scaling, the most significant difference between the ORS and \\iras\\ fields is induced by differing estimates of the over-density of the Virgo cluster. Neither of these gravity fields is consistent with the peculiar velocity field constructed from the full Mark III (\\cite{wil97a}) sample. We find that a simple but plausible non-linear bias algorithm for the galaxy distribution relative to the mass has a negligible effect on the derived fields. We conclude that the substitution of optical for \\iras\\ catalogues cannot alone resolve the discrepancies between the \\iras\\ gravity field and the Mark III peculiar velocity field. ", + "introduction": "With the advent of large, uniform redshift surveys of galaxies, increasingly stringent tests of models for the large-scale structure of the universe have become possible (see, e.g., \\cite{dek94} and \\cite{str95} for reviews). Under the assumption that large-scale flows of galaxies in the universe are a response to the underlying distribution of matter, peculiar velocity measurements are a critical probe of cosmology and large-scale structure. Redshift surveys allow one to construct an estimate of the gravity field based on the observed distribution of galaxies. Using linear or quasi-linear gravitational instability theory, one can predict the velocity field from this gravity field, a prediction which is approximately proportional to the parameter $\\beta \\equiv f(\\Omega)/b$, where $b$ is the galaxy bias, $f$ is a function which is well approximated by $\\Omega^{0.6}$ (\\cite{pee80}), and $\\Omega$ is the matter density. Until recently, the Infrared Astronomical Satellite (\\iras) 1.2-Jy flux-limited redshift survey was the only available nearly full-sky catalogue. The original redshift survey of \\iras-selected galaxies was limited to a sample of 2658 galaxies with 60-micron flux greater 1.936 Jy (\\cite{str92a}), but this was later extended to a sample of 5321 galaxies brighter than 1.2 Jy (\\cite{fis95a}). A sample of 13,000 galaxies complete to 0.6 Jy is nearing completion (\\cite{can98}) and is eagerly awaited. The principal advantage of these \\iras-selected samples is their insensitivity to extinction within the Milky Way, allowing extremely large sky coverage, $87.6$\\% (11.01 sr) for the 1.2-Jy sample, with a single, linear instrument. Nearly full-sky coverage is essential for estimation of the peculiar gravity field. However, it is well known that the distribution of \\iras\\ galaxies is biased with respect to the overall distribution of galaxies (\\cite{str92b}), primarily because \\iras\\ under-counts the dust-free early-type galaxies which congregate in cluster centers. The \\iras\\ surveys are also more dilute than is desirable, typically including only 1/3 of the known spiral galaxies in volumes where the selection function is high. Optically-selected surveys can eliminate this bias, and deeper redshift surveys will overcome the shot noise problem associated with the sparse sampling of the \\iras\\ survey. The recently completed Optical Redshift Survey (ORS; \\cite{san95}) is currently the best available catalogue for our peculiar velocity analysis. The ORS is a concatenation of three optically-selected samples covering most of the sky with $|b| > 20^\\circ$, and it is the best current approximation to a full-sky optically-selected catalogue. Until a redshift catalogue is constructed from the 2MASS survey (\\cite{ste95}), the ORS is likely to remain the closest competitor to the \\iras\\ full-sky catalogue. Previous work (\\cite{fre94}) suggests that the \\iras\\ and optical gravity fields are consistent, but with the superior data now available, it is clearly of interest to determine the nature of any deviations between the gravity fields of the optical and \\iras\\ samples, as these could substantially affect comparisons to observed peculiar velocity fields. Direct density field comparisons (\\cite{san92}; \\cite{str92b}) between the ORS and other redshift surveys will be carried out in detail by \\cite{str98}. Comparisons of the gravity field derived from the \\iras\\ survey with observed peculiar velocity samples have been carried out using a number of complementary methods (see \\cite{str95} for a review), with more recent studies performed by \\cite{dav96}, \\cite{wil97b}, \\cite{rie97}, \\cite{sig98}, \\cite{dac98}, and \\cite{wil98}. These studies show that the \\iras-predicted flow field is at least qualitatively very similar to the measured velocity field, lending strong support to the gravitational instability model for the growth of large-scale structure. Although the qualitative alignment of the \\iras-predicted and Mark III peculiar velocity fields is remarkably good, detailed quantitative comparisons are, at present, less satisfactory. DNW96 find large coherent residuals between the \\iras\\ and Mark III fields, primarily in the dipole component at large scales, which preclude a conclusive determination of $\\beta$. \\cite{wil97b} find better alignment of the \\iras\\ gravity and Mark III velocity fields using the rather different VELMOD technique, but this analysis is limited to redshifts within $3000\\kms$; \\cite{wil98} find good agreement on larger scales, but conclude that this requires a re-calibration of the Tully-Fisher relations for the various subsamples going into the Mark III catalogue inconsistent with the published calibration (\\cite{wil97a}). On the other hand, \\cite{dac98} find a better match between the large-scale \\iras\\ gravity field and the SFI I-band sample of spiral field galaxy peculiar velocity measurements (\\cite{gio97}). The SFI sample uses the same \\cite{mat92} data for the Southern sky as in Mark III, but the transformation of the data to a common system differs from the Mark III treatment. A recent summary of the current confusion is given by \\cite{dav98}. Might some of the discrepancy disappear if we were to use an optical redshift survey to construct an estimate of the gravity field? This paper addresses the question of the construction of a three-dimensional gravity field based on the ORS and \\iras\\ redshift catalogues. We construct our estimate of the gravity field using the redshift-space procedure described by \\cite{nus94}. Section 2 describes the galaxy catalogues and our procedure for merging them. The method for obtaining the gravity and velocity fields is briefly outlined in \\S 3, with details of the spherical harmonic formalism in the Appendix; more details have been presented elsewhere. Section 4 describes the differences between the ORS field and the field derived from \\iras\\ alone, including a brief discussion of alternatives to the simple linear biasing scheme, and \\S 5 discusses the main conclusions. ", + "conclusions": "The gravity field derived from the ORS is quite similar to the field derived from the \\iras\\ survey, provided the value of $\\beta$ is multiplied by a factor of $\\approx 2/3$, roughly consistent with the relative strength of clustering in the two surveys. Some differences from this simple scaling are observed; most significantly, the dipole of the ORS-predicted flow field is rotated towards Virgo and is further than the \\iras\\ prediction from the CMB dipole. This non-scaling behavior arises because the ratio of ORS to IRAS density fluctuations is larger in the local Supercluster than in other regions. DNW96 have found systematic discrepancies between the Mark III and \\iras-predicted flow fields, particularly a large coherent dipole residual, which cannot be physical because the dipole field depends only on the interior mass distribution. Estimating the gravity field from the ORS survey has provided a check on the influence of the known shortcomings of the \\iras\\ catalogue (namely, its sparse sampling and under-representation of cluster cores). Since the \\iras\\ and ORS gravity fields are so similar, we conclude that the redshift-space catalogues are not themselves the source of the DNW96 discrepancies. Compared to \\iras, the ORS gravity field will generally push us towards lower values of $\\beta$ in order to fit a given sample of peculiar velocity measurements. The result of DNW96 for the \\iras-Mark III comparison would then suggest $\\beta\\la 0.45$, but this comparison with Mark III remains suspect due to to the substantial quantitative discrepancies between the \\iras\\ and Mark III fields. \\cite{dac98}, on the other hand, showed that the SFI sample agrees quite well with the \\iras\\ gravity field for $\\beta=0.6$; presumably this sample would also be consistent with the ORS field for $\\beta\\approx 0.4$. Another serious concern for the Mark III comparison is the fact that complementary methods for attacking the velocity field problem have not converged to a consistent, unambiguous result. The POTENT analysis of \\cite{sig98} and \\cite{dek93} favors larger values ($\\beta\\sim 1$) than the ITF and VELMOD analyses. These methods weight the data in very different ways, and given that the fields show some level of inconsistency, it is not surprising that different analyses of the same data yield different answers. The sources of the DNW96 discrepancies between the \\iras\\ gravity field and the Mark III peculiar velocity measurements remain uncertain. These discrepancies might be telling us that galaxies cluster in a way which is not closely related to the clustering of the mass, or that the linear gravitational instability model provides an inadequate description of flows on the scales of interest. Alternatively, there may be unknown systematic biases in the Mark III data; see also the discussion in \\cite{wil98}. We have investigated a simple but plausible generalization of the usual linear bias model for galaxy clustering, but this does not significantly affect features in the predicted velocity field; even a rather sizeable change in the slope of the bias relation at $\\delta^g=0$ causes only a modest reduction in the amplitude of the field. Improved full-sky peculiar velocity surveys (e.g., \\cite{str97}) of larger galaxy samples are eagerly awaited. Such surveys are quite a difficult undertaking, but they may go some way towards clarifying our understanding of large-scale flows in the universe." + }, + "9802/astro-ph9802345_arXiv.txt": { + "abstract": "In proto-neutron stars with strong magnetic fields, the cross section for $\\nu_e$ ($\\bar\\nu_e$) absorption on neutrons (protons) depends on the local magnetic field strength due to the quantization of energy levels for the $e^-$ ($e^+$) produced in the final state. If the neutron star possesses an asymmetric magnetic field topology in the sense that the magnitude of magnetic field in the north pole is different from that in the south pole, then asymmetric neutrino emission may be generated. We calculate the absorption cross sections of $\\nue$ and $\\bnue$ in strong magnetic fields as a function of the neutrino energy. These cross sections exhibit oscillatory behaviors which occur because new Landau levels for the $e^-$ ($e^+$) become accessible as the neutrino energy increases. By evaluating the appropriately averaged neutrino opacities, we demonstrate that the change in the local neutrino flux due to the modified opacities is rather small. To generate appreciable kick velocity ($\\sim 300$~km~s$^{-1}$) to the newly-formed neutron star, the difference in the field strengths at the two opposite poles of the star must be at least $10^{16}$~G. We also consider the magnetic field effect on the spectral neutrino energy fluxes. The oscillatory features in the absorption opacities give rise to modulations in the emergent spectra of $\\nu_e$ and $\\bar\\nu_e$. ", + "introduction": "It has been recognized shortly after the discovery of pulsars that neutron stars have velocities much in excess of those of any other normal stellar populations in our Galaxy (Gunn \\& Ostriker 1970; Minkowski 1970). However, it is only in the last few years that a significant body of evidence has come into place to support the view that type II supernovae are nonspherical and neutron stars receive large kick velocities at birth. Lyne and Lorimer (1994) analyzed pulsar velocities in light of new proper-motion measurements (Harrison et al.~1993) and up-to-date pulsar distance scale (Taylor \\& Cordes 1993), and concluded that pulsars were born with a mean speed of $\\sim 450$~km~s$^{-1}$, much larger than previously thought. More recent studies of pulsar velocities have adopted more sophisticated satistical methods and included better treatment of selection effects and uncertainties (Lorimer et al.~1997; Hansen \\& Phinney 1997; Cordes \\& Chernoff 1997). They all yielded a result for the pulsar birth velocity in qualitative agreement (although not in quantitative details) with that of Lyne \\& Lorimer (1994). In particular, Cordes and Chernoff (1997) found that the three-dimensional space velocities of their sample of 47 pulsars have a bimodal distribution, with characteristic speeds of $180$~km~s$^{-1}$ and $700$~km~s$^{-1}$ (corresponding to $80\\%$ and $20\\%$ of the population). They also estimated that pulsars with velocities greater than $1000$~km~s$^{-1}$ may be underrepresented owing to selection effects in pulsar surveys [The uncertainty in the high-velocity end of the pulsar velocity distribution function has also be emphasized by Hansen \\& Phinney (1997)]. Concrete evidence for the existence of pulsars with velocities of $\\go 1000$~km~s$^{-1}$ has come from the observation of the Guitar Nebula pulsar (B2224+65), which produces a bow shock when plowing through the interstellar medium (Cordes, Romani \\& Lundgren 1993). In addition, studies of pulsar-supernova remnant associations have uncovered a number of pulsars having velocities much greater than $1000$~km~s$^{-1}$ (Frail et al.~1994), although in many cases the associations are not completely secure. Compelling evidence for supernova asymmetry and pulsar kicks also comes from the detection of geodetic precession in the binary pulsar PSR 1913+16 (Cordes et al.~1990; Arzoumanian et al.~1996), and the orbital plane precession in the PSR J0045-7319/B star binary (Lai et al.~1995; Kaspi et al.~1996) and its fast orbital decay (which indicates retrograde rotation of the B star with respect to the orbit; see Lai 1996). These results demonstrate that binary break-up (as originally suggested by Gott, Gunn \\& Ostriker 1970; see Iben \\& Tutukov 1996) can not be {\\it solely} responsible for the observed pulsar velocities, and that {\\it natal kicks are required}. In addition, evolutionary studies of neutron star binary population imply the existence of pulsar kicks (e.g., Deway \\& Cordes 1987; Fryer \\& Kalogera 1997; Fryer et al.~1998; see also Brandt \\& Podsiadlowski 1995). Finally, there are a large number of direct observations of nearby supernovae (e.g., Cropper et al.~1988; Trammell et al.~1993; McCray 1993; Utrobin et al.~1995) and supernova remnants (e.g., Morse, Winkler \\& Kirshner 1995; Aschenbach et al.~1995) in radio, optical, and X-ray bands which support the notion that supernova explosions are not spherically symmetric. The origin of the pulsar velocities is still unknown. Two classes of mechanisms for the {\\it natal kicks} have been suggested\\footnote{These exclude the slow, post-explosion rocket effect due to electromagnetic radiation from off-centered magnetic dipole moment (Harrison \\& Tademaru 1975).}. The first class relies on convective instabilities in the collapsed stellar core and within the rebounding shock (e.g., Burrows \\& Fryxell 1992; Burrows et al.~1995; Janka \\& M\\\"uller 1994,~1996; Herant et al.~1994). The asymmetries in the matter and temperature distributions associated with the instabilities naturally lead to asymmetric matter ejection and/or asymmetric neutrino emission. Numerical simulations indicate that the local, post-collapse instabilities are not adequate to account for kick velocities higher than a few hundred km~s$^{-1}$ (Janka \\& M\\\"uller 1994). A variant of this class of models therefore relies on the global asymmetric perturbations seeded in the presupernova cores (Goldreich et al.~1996; see also Bazan \\& Arnett 1994). Clearly, the magnitude of kick velocity depends on the degree of initial asymmetry in the imploding core (Burrows \\& Hayes 1996). Due to various uncertainties in the presupernova stellar models, it is not clear at this point whether sufficiently large initial perturbations can be produced in the precollapse core (Lai \\& Goldreich 1998). In this paper we focus on the second class of models in which the large pulsar kick velocities arise from asymmetric neutrino emission induced by strong magnetic fields. Since $99\\%$ of the neutron star binding energy (a few times $10^{53}$~erg) is released in neutrinos, tapping the neutrino energy would appear to be an efficient means to kick the newly-formed neutron star. Magnetic fields are naturally invoked to break the spherical symmetry in neutrino emission. But the actual mechanism is unclear. We first briefly comment on previous works on this and related subjects. \\subsection{Previous Works on Asymmetric Neutrino Emission Induced by Magnetic Fields} A number of authors have noted that parity violation in weak interactions may lead to asymmetric neutrino emission from proto-neutron stars (Chugai 1984; Dorofeev et al.~1985; Vilenkin 1995; Horowitz \\& Piekarewicz 1997). However, their studies are largely unsatisfactory for a number of reasons: they either failed to identify the most relevant neutrino emission/interaction processes or the relevant physical conditions in proto-neutron stars, or stopped at estimating the magnetic field effects on neutrino opacities. Chugai (1984) and Vilenkin (1995) (see also Bezchastnov \\& Haensel 1996) considered neutrino-electron scattering and concluded that the effect is extremely small\\footnote{Note that Chugai's estimate for the electron polarization in the relativistic and degenerate regime (the relevant physical regime) is incorrect. This error leads to an overestimate of the effect as compared to Vilenkin's result.} (e.g., to obtain $V_{\\rm kick}=300$~km~s$^{-1}$ would require a magnetic field of at least $10^{16}$~G). However, neutrino-electron scattering is less important than neutrino-nucleon scattering in determining the characteristics of neutrino transport in proto-neutron stars. Similarly, Dorofeev et al.~(1985) considered neutrino emission by Urca processes in strong magnetic fields. But as shown by Lai \\& Qian (1998), in the bulk interior of the neutron star, the asymmetry associated with neutrino emission is cancelled by the asymmetry associated with neutrino absorption. Concerning the parity violation effect in proto-neutron stars, Horowitz \\& Li (1997) recently suggested that the asymmetry in neutrino emission may be enhanced due to multiple scatterings of neutrinos by nucleons which are slightly polarized by the magnetic field. They estimated that a field strength of a few times $10^{12}$~G is adequate to account for kick velocities of a few hundred km~s$^{-1}$. However, the paper by Horowitz \\& Li (1997) only discusses the idealized situations of scattering media, and ignores various essential physics that is needed in a proper treatment of neutrino transport in proto-neutron stars. In particular, it does not consider the effect of neutrino absorption which one might suspect to wash out the cumulative effect from multiple scatterings. Our own study of the parity violation effect (Lai \\& Qian 1998) was flawed in the treatment of scattering terms in the neutrino transport equation. As shown by Arras \\& Lai (1998), detailed balance requires that there be no cumulative effect from multiple scatterings in the bulk interior of the proto-neutron star where local thermodynamic equilibrium applies to a good approximation. Enhancement of neutrino emission asymmetry from multiple scatterings obtains only after neutrinos thermally decouple from proto-neutron star matter, and therefore is insignificant. Bisnovatyi-Kogan (1993) attributed pulsar kicks to asymmetric magnetic field distribution in proto-neutron stars. Using the fact that neutron decay rate can be modified by the magnetic field, he inferred that neutrino emissions from opposite sides of the neutron star surface are different. However, neutron decay is not directly relevant for neutrino emission from a newly-formed neutron star. Roulet (1997) (whose paper was posted while this paper was being written) considered the relevant neutrino absorption processes. But his study was restricted to calculating the neutrino cross sections as functions of neutrino energy (corresponding to \\S2.1 of this paper), and therefore was insufficient for addressing the issue of asymmetric neutrino emission. Finally, it is worthwhile to mention a speculative idea on pulsar kicks which relies on nonstandard neutrino physics. Kusenko \\& Segr\\`e (1996) suggested that asymmetric $\\nu_\\tau$ emission could result from the Mikheyev-Smirnov-Wolfenstein flavor transformation between $\\nu_\\tau$ and $\\nu_e$ inside a magnetized proto-neutron star because a magnetic field changes the resonance condition for the flavor transformation. With a $\\nu_\\tau$ mass of $\\sim 100$~eV, they claimed that magnetic fields of $\\sim 3\\times 10^{14}$~G can give the pulsar a kick velocity of a few hundred km~s$^{-1}$. However, their treatment of neutrino transport was oversimplified: e.g., they ignored that neutrinos of different energies have different resonance surfaces. Furthermore, a simple geometric effect and realistic proto-neutron star conditions easily reduce their estimated kick velocity by an order of magnitude. Most likely, a magnetic field strength of $\\sim 10^{16}$~G is needed to produce the observed average pulsar kick velocity via their mechanism (see Qian 1997). Similarly strong magnetic fields are required in variants of the Kusenko \\& Segr\\`e mechanism, e.g., those considered by Akhmedov et al.~(1997) (which relies on both the neutrino mass and the neutrino magnetic moment to facilitate the flavor transformation). \\subsection{Plan of This Paper} As should be clear from the brief review in \\S1.1, in spite of quite a number of suggestions, there is at present no consensus on the magnitudes of the magnetic field induced asymmetry in neutrino emission from proto-neutron stars and the resulting kick velocities. In this paper we study the effect of asymmetric magnetic field topology on pulsar kicks. Since the energy levels of $e^-$ and $e^+$ in a magnetic field are quantized, the $\\nu_e$ and $\\bnue$ absorption opacities near the neutrinosphere and the neutrino-matter decoupling region depend on the local magnetic field strength. If the {\\it magnitude} of magnetic field in the north pole is different from that in the south pole (the field does not need to be ordered), then asymmetric neutrino flux may be generated. Here we ignore the effects of magnetic fields on the proto-neutron star structure through the equation of state --- these are secondary effects as far as neutrino transport is concerned. The organization of this paper is as follows. In \\S2 we calculate the $\\nu_e$ and $\\bar\\nu_e$ absorption cross sections as functions of neutrino energy in strong magnetic fields. Section 3 summarizes the basic features of neutrino transport near the neutron star surface. In \\S4 we derive and evaluate the ``Rosseland mean'' neutrino opacities which are directly related to the local neutrino flux. We then consider in \\S5 the change in neutrino flux due to the modified opacities, and estimate the kick velocity resulting from an asymmetric magnetic field topology. In \\S6 we briefly consider how the emergent neutrino spectra may be modified by the strong magnetic field. In \\S7 we present our conclusion. Unless noted otherwise, we shall use units in which $\\hbar$, $c$, and the Boltzmann constant $k$ are unity. ", + "conclusions": "In this paper, we have studied the issue of whether asymmetric magnetic field topology in a proto-neutron star can induce asymmetric neutrino emission from the star through the modifications of the neutrino absorption opacities by the magnetic field. These modifications arise from the quantized Landau levels of electrons and positrons produced in strong magnetic fields. By calculating the appropriate mean neutrino opacities in the magnetic field, we demonstrate that this mechanism is rather inefficient in generating a kick to the neutron star: To obtain appreciable kick velocities ($\\sim 300$~km~s$^{-1}$), the difference in the field strengths at the two opposite poles of the star must be at least $10^{16}$~G." + }, + "9802/hep-ph9802393_arXiv.txt": { + "abstract": "The effect of multibody massless neutrino exchanges between neutrons inside a finite-size neutron star is studied. We use an effective Lagrangian, which incorporates the effect of the neutrons on the neutrinos. Following Schwinger, it is shown that the total interaction energy density is computed by comparing the zero point energy of the neutrino sea with and without the star. It has already been shown that in an infinite-size star the total energy due to neutrino exchange vanishes exactly. The opposite claim that massless neutrino exchange would produce a huge energy is due to an improper summation of an infrared-divergent quantity. The same vanishing of the total energy has been proved exactly in the case of a finite star in a one-dimensional toy model. Here we study the three-dimensional case. We first consider the effect of a sharp star border, assumed to be a plane. We find that there is a non-vanishing of the zero point energy density difference between the inside and the outside due to the refraction index at the border and the consequent non-penetrating waves. An analytical and numerical calculation for the case of a spherical star with a sharp border confirms that the preceding border effect is the dominant one. The total result is shown to be infrared-safe, thus confirming that there is no need to assume a neutrino mass. The ultraviolet cut-offs, which correspond in some sense to the matching of the effective theory with the exact one, are discussed. Finally the energy due to long distance neutrino exchange is of the order of $10^{-8}$--$10^{-13}\\ {\\mathrm{GeV}\\hbox{ per neutron}}$, i.e. negligible with respect to the neutron mass density. ", + "introduction": "The massless neutrino exchange interaction between neutrons, protons, etc., is a long-range force \\cite{Fein68}--\\cite{sikivi}. In a previous work \\cite{rescue}, the long-range interaction effects on the stability of a neutron star due to multibody exchange of massless neutrinos have been studied. We have shown that the total effect of the many-body forces of this type results in an infrared well-behaved contribution to the energy density of the star and that it is negligible with respect to the star mass density. This is in agreement with two recent non-perturbative calculations done by Kachelriess \\cite{Kachelriess} and by Kiers and Tytgat \\cite{kiers}. This work is in contradiction with the repeated claim by Fischbach \\cite{fischbach} that, unless the neutrino is massive, neutrino exchange renders a neutron star unstable, as the induced self-energy exceeds the mass of the star because of the infrared effects associated to neutrino exchange between four or more neutrons. In our opinion, the latter ``catastrophic'' result is a consequence of summing up large infrared terms in perturbation outside the radius of convergence of the perturbative series. The non-perturbative use of an effective Lagrangian immediately gives the result without recourse to the perturbative series, and the result is small. Smirnov and Vissani \\cite{smirnov}, following the same method as in ref. \\cite{fischbach}, summing up multibody exchange contributions order by order, showed that the 2-body contribution is damped by the blocking effect of the neutrino sea \\cite{loeb}. They guessed that this damping would apply to many-body contributions, and hence would reduce the catastrophic effect claimed by Fischbach. In our previous work \\cite{rescue}, we also considered the effect of the neutrino sea inside the neutron star. This effect has been introduced in our non-perturbative calculation by using Feynman propagators of neutrinos inside a dense medium, which incorporate the condensate term. We noticed that this condensate is present, but in our opinion it is not the most important of the effects that neutralize the catastrophic effect expected by Fischbach, since it only brings a tiny change to the non-perturbatively summed interaction energy, leaving unchanged our conclusion that the weak self-energy is infrared-safe. We have also stressed in \\cite{rescue} that the neutrino condensate was related to the existence of a border \\footnote{ At that point, we should call the reader's attention to a minor mistake that was made in the previous calculation \\cite{rescue}: a pole was forgotten in the calculation of the weak self-energy, and its contribution is exactly annihilated by the condensate's, as shown in ref. \\cite{note}: the computation of the self-energy gives zero when the forgotten pole is taken into account.}. This was demonstrated in the ($1\\ +\\ 1$)-dimensional star in \\cite{note}: the blocking effect, which implies the trapping of the neutrinos inside the star while the antineutrinos are repelled from it, is a natural consequence of the existence of a border. Indeed a proper treatment of the effect of the border automatically incorporates the condensate contribution as a consequence of the appropriate boundary conditions for the neutrino Feynman propagator inside the star. Since our treatment directly incorporates the neutrino sea effect and, on the other hand, since we stick to our strategy of directly computing the total neutrino interaction energy by using an effective Lagrangian, we have in a sense thus generalized the result of Smirnov and Vissani \\cite{smirnov} because our result holds in a non-perturbative way and accounts for all the $n$-body contributions, while theirs holds for $2$-body contributions only. It has been objected \\cite{smirnov1} to our study \\cite{rescue} that we worked in an approximation where we neglected the border of the neutron star. Our belief is that this simplifying hypothesis does not change the fundamental result that the total effect of the multiple neutrino exchange to the energy density of the star is not infrared-divergent. Indeed, in \\cite{note} we have proved that, in $(1\\ +\\ 1)$ dimensions {\\it with borders}, the result of \\cite{rescue} for the infinite star {\\it without borders} is kept unchanged: the net interaction energy due to long-range neutrino exchange is exactly zero. This result was obtained by computing Feynman vacuum loops with neutrino propagators derived from the effective Lagrangian, which incorporates the neutron interaction. The latter propagators are not translational-invariant, because of the star borders. We have found a physically simple explanation for the vanishing of the net interaction energy. It relies on the fact that the negative energy states, in the presence of the star and without the star, are in a one-to-one correspondence and have exactly the same energy density. It results that the zero point energy is the same with and without the star, for any density profile of the star. The main goal of the present work is to follow on taking into account the finite-size and border effects. Two main conclusions of ref. \\cite{note} are useful for the ($3\\ +\\ 1$)-dimensional star: i) the natural connection between the neutrino sea and the border, ii) the correct definition of the zero-energy level of the Dirac sea. From \\cite{note}, we know that the zero-energy level of the Dirac sea has to be adjusted by comparing the asymptotic behaviour of the wave functions far outside the star with the free solutions in the absence of the star. From there, we know the correct $i \\epsilon$ prescription, which has to be imposed in the propagators of the neutrinos in the presence of the star; we could in principle compute the closed loops to get the vacuum energy density. However, this approach is technically very difficult. A simpler method, the derivation of which is recalled in section \\ref{general}, is to simply add the energy density of the negative energy solutions in the presence of the star, minus the same in the absence of the star. The vanishing of the neutrino exchange energy density in the star, found in the case of an infinite star \\cite{rescue} and in $(1\\ +\\ 1)$ dimensions \\cite{note}, which will be summarized in section \\ref{1+1}, is not valid in $(3\\ +\\ 1)$ dimensions. The main reason for that will be illustrated in section \\ref{flat} by zooming to the border effect, i.e. considering a plane border in $(3\\ + \\ 1)$ dimensions. There is a non-trivial refraction index that modifies the wave energy densities as they penetrate the star. Some waves are forbidden to penetrate and this yelds the dominant contribution. In section \\ref{spherical}, we perform analytical and numerical calculations, which take into account the curvature of the border and use a spherical star with a sharp border. We find a very simple approximate formula for the zero point neutrino energy density in the star, and demonstrate numerically the validity of this approximation. It results that indeed the neutrino-induced energy density in the star does not vanish and is dominantly explained by the above-mentioned border effect. In any case, these non-vanishing neutrino exchange energy densities are all perfectly regular in the infrared and do not present any resemblance to Fischbach's effect. On the other hand they are ultraviolet-singular. This is not unexpected since anyhow our effective Lagrangian is only valid below some energy scale where the neutrons may be considered at rest. In section \\ref{discussion}, we also discuss in some detail the effect of decoherence for distances larger than the neutrino mean-free path, which smoothes down the ultraviolet singularity. We conclude that the stability of compact and dense objects such as a white dwarfs, neutron stars, etc., are not affected by the neutrino exchange, even if neutrinos are massless. ", + "conclusions": "\\label{remarks} In order to settle definitively the question of the stability of a neutron star, the multibody exchange of massless neutrinos has been computed analytically and numerically for a finite star. The effect of a border is twofold. First it induces in a natural way the neutrino condensate as proved in \\cite{note}. The latter condensate, does not produce any neutrino exchange interaction energy in the simplified $(1\\ +\\ 1)$-dimensional case and we find it to be negligible in the realistic $(3\\ +\\ 1)$-dimensional case. The second effect of the border is that the neutrino zero point energy inside the star differs from the outer one because of negative-energy waves that cannot penetrate inside the star, being beyond the limiting refraction index at the border. This contribution is proportional to the volume of the star, but it is still tiny ($10^{-8}$--$10^{-13} \\ {\\mathrm{GeV}\\hbox{ per neutron}}$), completely negligible in comparison with the neutron mass. {\\it We find no infrared divergences in the full non-perturbative result, which would have necessitated the introduction of a neutrino mass.} The general conclusion of this work is that the neutrino does not need to be massive to ensure the stability of a neutron star. This is in agreement with recent works (commented below) by Kachelriess \\cite{Kachelriess} and by Kiers \\& Tytgat \\cite{kiers}. There is no catastrophic effect due to the multibody massless neutrino exchange. As already stated in refs. \\cite{rescue} and \\cite{note}, this catastrophic result claimed in \\cite{fischbach} is only due to an attempt to sum up the perturbative series outside its radius of convergence. While finishing this paper, there appeared a paper by Kiers and Tytgat \\cite{kiers}. They accept the point of view developed in \\cite{rescue} for an infinite star and try, as we do, to solve the problem of a finite star. Their starting goal is, as ours, to compute the density in eq. (\\ref{somme}). They use a clever technique based on quantizing in a large sphere and expressing the vacuum energy density in terms of the phase shifts. They first study analytically the unphysical but illustrative case of small $bR$ and then numerically the large $bR$ case. They show that the perturbative series {\\it \\`a la Fischbach} already breaks down as early as $b R> \\pi$ (for the neutron star, $b R\\sim 10^{12}$) while {\\it the non-perturbative calculation gives a negligible result, a conclusion which we fully share}. One difference between their result and ours is that they find a relation between the energy density of the neutron star and that of the neutrino condensate. We find on the contrary that the result is mainly due to non-penetrating negative-energy waves, which are not related to the condensate. We did not yet manage to understand the reason for this discrepancy. It might be related to different UV-regularization methods. Kachelriess \\cite{Kachelriess} also agrees with us about Fischbach's ``catastrophic'' result. He computed the total weak self-energy for an infinite neutron star following the Schwinger method, by using the neutrino propagator in momentum space, as we did in ref. \\cite{rescue}. He obtained a non-zero weak self-energy without taking into account the neutrino sea effects. As we acknowledged in section \\ref{flat} and in ref. \\cite{note}, a minor mistake was made in ref. \\cite{rescue}: the contribution of a pole had been forgotten in the calculation. Once this mistake is corrected, we agree with Kachelriess about the result when neglecting the neutrino sea. He attributed the discrepancy to the fact that we took the limit ($y \\to x$) before integrating over the whole space to obtain the total weak self-energy (see eq. (\\ref{Wanalyt})). We do not agree with this conjecture about the discrepancy, first because, as we just mentioned, the forgotten pole removes the discrepancy, second because we have verified that our previous UV-regularization makes the $x\\to y$ limit regular. Finally, from our analysis of finite stars \\cite{note}, we insist that the condensate {\\it has to be taken into account} and, amazingly, it exactly cancels the forgotten pole contribution, resulting finally in $w(\\vec x)=0$ for an infinite star. A few weeks later, Fischbach and Woodahl \\cite{repeated} repeated Fischbach's original claim and used the same expansion, order by order, in the number of neutrons, i.e. equivalently in perturbation in the parameter $bR$. Astonishingly enough they did not consider the series of works demonstrating that this series is simply divergent, but that the total result may be computed directly by the effective Lagrangian technique. They argued that our non-perturbative calculation encounters cancellations because, in the effective approach, the neutron medium is assumed by us to be a continuous background. Of course, treating the neutron medium as a homogeneous continuum medium is an approximation \\`a la Hartree--Fock, and it should be corrected {\\it in the ultraviolet} by taking into account the correlation between neutrons. This is precisely one of the reasons why we considered that a natural ultraviolet cut-off was the energy scale of a few MeV. The authors of \\cite{fischbach} take the size of the neutron hard core as an ultraviolet cut-off. Why not, although many other ultraviolet effects arise at the same scale of a few 100 MeV: the neutron recoil, the quark and gluon content of the neutron, without forgetting the incoherent neutrino--neutron scattering discussed in the previous section. But these ultraviolet effects {\\it will not at all modify the analysis of the infrared catastrophe} advocated by Fischbach and denied by us. The authors of \\cite{repeated} seem to imply that we have added some unjustified assumption in our work. The truth is on the contrary that we have assumed nothing that they did not assume themselves, such as the static neutron assumption, but we have not assumed, as they do, that a neutron is not allowed to interact more than once, neither did we make the drastic approximations that appear in their work at high order in perturbation. {\\it The effective Lagrangian approach allows to compute exactly, in a simple manner, and with fewer assumptions than the perturbative expansion approach.} The authors of \\cite{repeated} criticize our recent $(1\\ +\\ 1)$-dimensional toy calculation \\cite{note} arguing that the critical parameter $bR$ is much smaller than 1 in $(1\\ +\\ 1)$ dimensions. However, they did not notice that our $(1\\ +\\ 1)$-dimensional result is absolutely exact, independently of the parameter $bR$, which, incidentally, we have taken to be large. To finish, we feel it necessary to insist. The main issue is the failure of the perturbative expansion, which is infrared-divergent. Happily one can spare this difficulty thanks to the effective action technique. Once this point is understood, the different analyses all agree, notwithstanding minor discrepancies, that {\\it although the massless neutrino exchange between fermions is a long-range interaction, it does not give any significant contribution to the total energy of a neutron star, finite or infinite}." + }, + "9802/astro-ph9802351_arXiv.txt": { + "abstract": "The complete set of data from the Tenerife 10 GHz ($8\\dg$ FWHM) twin-horn, drift scan experiment is described. These data are affected by both long-term atmospheric baseline drifts and short term noise. A new maximum entropy procedure, utilising the time invariance and spatial continuity of the astronomical signal, is used to achieve a clean separation of these effects from the astronomical signal, and to deconvolve the effects of the beam-switching. We use a fully positive/negative algorithm to produce two-dimensional maps of the intrinsic sky fluctuations. Known discrete sources and Galactic features are identified in the deconvolved map. The data from the 10 GHz experiment, after baseline subtraction with MEM, is then analysed using conventional techniques and new constraints on Galactic emission are made. ", + "introduction": "Anisotropy of the cosmic microwave background (CMB) radiation provides one of the key constraints against which cosmological models can be tested. However, the detection of anisotropy at a level $\\Delta T/T$ $\\sim 10^{-5}$ is a challenging observational problem. Attainment of these sensitivity levels is only possible by making differential measurements, switching between two different sky patches or by using interferometric techniques. The data from the telescope consist of the true sky brightness distribution convolved in some instrument beam, with an additional random noise contribution. Combined with the often non-uniform scan strategies of CMB experiments this makes the task of deconvolution to obtain a two-dimensional map of the intrinsic fluctuations a non-trivial proposition. Here we describe the analysis procedures that have been developed and applied to the Tenerife drift scan data (Davies \\et 1987\\nocite{davies87}, Watson \\et 1992\\nocite{watson92}, Hancock \\et 1994\\nocite{h94}, Guti\\'errez \\et 1997) in order to produce 2-D maps of the intrinsic fluctuations, while at the same time providing scans free from atmospheric baseline variations. Although the implementation described here is specific to the Tenerife instruments, some of the techniques are more generally applicable to other CMB data sets (see e.g. Maisinger {\\em et al} 1997) and offer a useful means of comparing between observations from telescopes with different beam patterns and scan strategies. In order to demonstrate the technique we consider here the analysis of the total data set from the original (FWHM$\\sim 8\\dg$) Tenerife twin horn radiometer experiment (Davies \\et 1992\\nocite{davies92}). Although the instrument configuration differs from that of the current telescopes (Davies \\et 1996a\\nocite{davies95}) and is less sensitive to cosmological signals (Watson \\et 1992\\nocite{watson92}), this experiment has surveyed a substantial fraction of the full sky, making it interesting to attempt a 2-D mapping. While a partial analysis of a limited RA range, along a strip at Dec $= +40\\dg$ has been given elsewhere (Davies \\et 1987\\nocite{davies87}), we here present a thorough analysis of all of the data, covering a selection of declinations ranging from $-15\\dg$ to $+45\\dg$. At the operating frequency of 10.4 GHz our maps will be sensitive to synchrotron and free-free emission in addition to the CMB and discrete radio sources. These maps can be used to constrain contaminating signals in higher frequency observations by virtue of the differing spectral dependencies of the CMB and the foregrounds. Testing the procedures on known structures will also give us improved confidence when analysing the higher frequency 15 and 33 GHz data to obtain 2-D maps of the CMB. In Section 2 we present an analysis scheme utilising a maximum entropy based regularising function, to enable a clean separation of the astronomy from spurious atmospheric effects and to provide a deconvolved 2-D image of the microwave sky. In Section 3 we present details of the observations and the implementation of the algorithm. In Section 4 we test the positive/negative maximum entropy method with simulations of the observations. Section 5 details the application of the algorithm to the data from the Tenerife experiment. Sections 6, 7 and 8 describe the separation of cosmological and astronomical signals, the application of suitable significance tests and the interpretation and conclusions that were reached. ", + "conclusions": "We have presented here a new method for analysing the data from microwave background experiments. As seen from simulations performed in Section \\ref{sim}, the positive/negative MEM algorithm performs very well recovering the amplitude, position and morphology of structures in both the reconvolved scans and the two-dimensional deconvolved sky map. We conclude that no bias, other than the `damping' enforcement, is introduced into the results from the methods described here, as the bare minimum of prior knowledge of the sky is required. A simultaneous baseline fit is also possible. Even with the lowest signal to noise ratio (the $100\\mu$K noise simulation which corresponds to our worse case in the Tenerife experiments) all of the main features on the sky were reconstructed. Using this method we are able to put constraints on the galactic contamination for other experiments at higher frequencies which is essential when trying to determine the level of CMB fluctuations present. It is clear that this approach works well and provides a useful technique for extracting the optimum CMB maps from both current and future multi-frequency experiments. This will become of ever increasing importance as the quality of CMB experiments improves. At present we are using this method to analyse the new data from the three beam switching Tenerife experiments at 10 GHz, 15 GHz and 33 GHz with angular resolutions of $\\sim 5\\degg$. We hope the two dimensional results from these will be available shortly. \\subsection*{ACKNOWLEDGEMENTS} \\noindent The Tenerife experiments are supported by the UK Particle Physics and Astronomy Research Council, the European Community Science programme contract SCI-ST920830, the Human Capital and Mobility contract CHRXCT920079 and the Spanish DGICYT science programme. A.W. Jones wishes to acknowledge a UK Particle Physics and Astronomy Research Council Studentship. S. Hancock wishes to acknowledge a Research Fellowship at St. John's College, Cambridge, UK. G. Rocha wishes to acknowledge a JNICT Studentship from Portugal." + }, + "9802/astro-ph9802167_arXiv.txt": { + "abstract": "\\noindent The current formation models for cluster elliptical galaxies which incorporate a mechanism for the metallicity enhancement of massive ellipticals predict a change in the observed slope of the red sequence of ellipticals as a function of redshift. This change occurs primarily because the metal-rich galaxies become redder faster than the metal-poorer galaxies with increasing age. This effect is most pronounced within $\\sim$ 4 Gyr of formation. Observations of the change of the slope of the red sequence with redshift may thus be used to constrain the formation epoch for galaxy clusters. We examine the red sequence of cluster ellipticals using publicly available HST imaging data for a set of six $0.75>$ z $>0.2$ clusters, and a sample of 44 Abell clusters at z $<0.15$ imaged with the KPNO 0.9 m. We compare the derived slopes of the red sequences with a set of cluster-elliptical evolution models and find good agreement. We demonstrate that such a comparison provides a useful constraint on the formation epoch for clusters, which can be made independently from considerations of absolute color evolution and scatter in the red sequence. From our initial comparison of the observed and model slopes as a function of redshift, we conclude, conservatively, that most of the elliptical galaxies in the cores of clusters must form at z $>2.0$, and that these galaxies are coeval and passively evolving. ", + "introduction": "\\noindent The existence of a color-magnitude relation (CMR) for field ellipticals, in which the brighter ellipticals are in general redder, was first noted by Baum (1959). Locally, the elliptical galaxies in individual clusters form a red sequence, with a well--defined slope and small scatter (Bower\\etal 1992a, b). Recent results from the Hubble Space Telescope demonstrate the existence of a tight red sequence, comparable in scatter to that observed in the red sequence of the Coma cluster, in clusters at redshifts up to $z=0.9$ (Stanford\\etal 1998, hereafter SED98; Ellis\\etal 1997, hereafter E97). Though the CMR in clusters can be interpreted at the present epoch as either an age effect or a metallicity effect, the existence of the red sequence at redshifts greater than 0.3 makes the age explanation untenable (Kodama 1997, hereafter K97). Specifically, the existence of the red sequence at higher redshifts indicates that cluster ellipticals are a passively evolving population in which the reddening of massive galaxies is the result of a mass-metallicity relation rather than an age effect (Kodama \\& Arimoto 1997; Kauffman \\& Charlot 1997; K97). \\paragraph{} The apparent passive evolution of cluster ellipticals is broadly consistent with models in which ellipticals form in a monolithic collapse at high redshift, and evolve passively after this initial star-burst (Eggen\\etal 1962). The origin of the mass-metallicity relation in this collapse scenario was first explored in detail by Arimoto \\& Yoshii (1987), who included the effects of supernova winds (Larson 1974). The heating of the interstellar medium by supernovae in the initial star-burst triggers the formation of a galactic wind when the thermal energy of the gas exceeds the gravitational binding energy. This galactic wind ejects the gas in low-mass galaxies more efficiently due to their shallower potential wells, resulting in a trend of increasing metallicity with mass (e.g., Carlberg 1984a, b). The more massive galaxies are more likely to retain the enriched supernova ejecta, and have star-bursts of longer duration. \\paragraph{} The initial enrichment difference between elliptical galaxies of different masses is manifested in the present epoch as the slope of the red sequence, which appears to be constant from cluster to cluster (Lopez-Cruz \\& Yee 1998; Lopez-Cruz 1997, hereafter LC97). However, this mass-metallicity relation causes elliptical galaxies of different masses to display slightly differing photometric properties with age. This differing color evolution manifests as a change in the slope of the red sequence with redshift. In a scenario in which elliptical galaxies in the cores of clusters form concurrently in a monolithic collapse at high redshift, we expect the observed slope of the red sequence to be significantly flatter at ages less than $\\sim$ 4 Gyr after formation (K97; Kodama \\& Arimoto 1997), when compared to later times. \\paragraph{} An alternative explanation for the origin of the slope of the red sequence is provided by Kauffman \\& Charlot (1997), in which elliptical galaxies are formed hierarchically through the merger of disk systems. In this model there is again a mass-metallicity relation for the progenitor disk systems, and the enrichment occurs prior to the assembly of elliptical galaxies. Massive elliptical galaxies tend to form from the more massive disk systems in hierarchical merging, resulting in a mass-metallicity relation for elliptical galaxies. In this model the observed slope of the red sequence is also expected to flatten at high redshifts, because the stellar populations in massive ellipticals are on average younger, and become bluer relative to low-mass systems as the formation epoch is approached. \\paragraph{} In theory, then, the change in the slope of the red sequence with redshift can be used to constrain the formation epoch of the dominant stellar population in elliptical galaxies (K97; LC97). This constraint can be made independently of the intrinsic scatter in early-type galaxy colors about the red sequence (Bower\\etal 1992a,b; Stanford\\etal 1995; E97; SED98), and the evolution of the apparent color with redshift (Arag\\'{o}n-Salamanca\\etal 1993; K97; E97; SED98; Kodama\\etal 1997). The measurement of the slope offers a significant advantage over measurements of color and scatter in that it is calibration independent. The slope can be measured with the same precision even in the presence of large systematic errors in photometric zero points. \\paragraph{} In this paper we present an analysis of the slopes of the red sequences for a total of 50 clusters spanning the redshift range 0$<$ z $<$0.75, and demonstrate that such an analysis offers a powerful tool in constraining the formation epoch of elliptical galaxies in clusters. Our data are drawn from two sources: a subset of the imaging survey of LC97 of 45 Abell clusters at z $<$0.2, and six clusters at z $>$0.2 drawn from the Hubble Space Telescope (HST) archive. We consider only the interior 0.5 \\Mpc of each cluster, and select early-type galaxies in the high redshift sample using an automated galaxy classifier based on central concentration (c.f. Abraham\\etal 1994). In \\S 2, we define the imaging dataset in detail, and discuss the galaxy photometry and morphology parameters derived from these data. In \\S 3, we present our analysis of the red sequence slopes in these clusters, and compare these results to the models of K97 in \\S 4. In \\S 5 we summarize our conclusions, and comment on the applicability of this analysis to even higher redshift observations. We use $H_{0}$=50, $\\Omega_{{\\it m}}$=1.0, and $\\Omega_{\\Lambda}$=0.0 throughout this paper, unless otherwise noted. ", + "conclusions": "" + }, + "9802/gr-qc9802066_arXiv.txt": { + "abstract": "We present a method for finding the eigenmodes of the Laplace operator acting on any compact manifold. The procedure can be used to simulate cosmic microwave background fluctuations in multi-connected cosmological models. Other applications include studies of chaotic mixing and quantum chaos. ", + "introduction": "Much of physics boils down to solving differential equations subject to certain boundary conditions. Waves in a box are a classic example. The box we have in mind here is some closed manifold, and the waves are those of a massless scalar field. This picture arises naturally when studying density perturbations in a multi-connected universe\\cite{ll}. The mathematical problem can be stated: Find all square integrable functions $\\Psi_q({\\bf x})$ that satisfy the partial differential equation \\begin{equation} (\\nabla + q^2 )\\Psi_q({\\bf x}) = 0 \\, . \\end{equation} Here $\\nabla$ is the Laplace operator on some closed manifold $\\Sigma$ and the constant $q$ is an eigenvalue of the Laplacian. The complexity of the problem is controlled by the geometry of $\\Sigma$. When $\\Sigma$ is $n$-dimensional Euclidean space, $E^n$, modulo some discrete group of covering transformations, $\\Gamma$, it is a simple matter to write down analytic expressions for the eigenmodes $\\Psi_q({\\bf x})$. In contrast, if $\\Sigma$ describes some compact hyperbolic manifold, $H^n / \\Gamma$, then the eigenmodes cannot be expressed in closed analytic form. This difficulty is closely related to the chaotic behaviour of geodesic flows on compact negatively curved spaces\\cite{berry}. Here we describe a numerical solution to the problem based on Fourier filtering solutions to the scalar wave equation $\\Box \\Psi =0$. While our approach works for any topology, we will focus on hyperbolic manifolds as these are of the most interest to cosmology. ", + "conclusions": "Having demonstrated that our approach works in 2-dimensions, our next task is to apply it to the cosmologically relevant case where $\\Sigma$ is a compact hyperbolic 3-manifold. The computational cost of having an extra dimension will be offset by the availability of examples with very small fundamental domains. Many small hyperbolic 3-manifolds have FD's that have outradii, $\\eta_+$, smaller than the radius of curvature\\cite{jeff}. Consequently, space looks approximately Euclidean inside the FD, so standard coordinate systems such as the Klein metric will produce fairly uniform computational grids. Once we have found all the low lying eigenmodes we can use them to produce simulated maps of the cosmic microwave background radiation. These maps can then be compared to observational data, or used to test proposals for finding the large scale topology of the universe\\cite{cssj}." + }, + "9802/astro-ph9802021_arXiv.txt": { + "abstract": "The global geometry of the universe is in principle as observable an attribute as local curvature. Previous studies have established that if the universe is wrapped into a flat hypertorus, the simplest compact space, then the fundamental domain must be at least $0.4$ times the diameter of the observable universe. Despite a standard lore that the other five compact, orientable flat spaces are more weakly constrained, we find the same bound holds for all. Our analysis provides the first limits on compact cosmologies built from the identifications of hexagonal prisms. ", + "introduction": " ", + "conclusions": "" + }, + "9802/astro-ph9802092_arXiv.txt": { + "abstract": "We combine the photometric redshift data of Fern\\'andez-Soto {\\it et al.} (1997) with the morphological data of Odewahn {\\it et al.} (1996) for all galaxies with $I < 26.0$ detected in the Hubble Deep Field. From this combined catalog we generate the morphological galaxy number-counts and corresponding redshift distributions and compare these to the predictions of {\\it high normalization} zero- and passive- evolution models. From this comparison we conclude the following: \\noindent (1) E/S0s are seen in numbers and over a redshift range consistent with zero- or minimal passive- evolution to $I = 24$. Beyond this limit fewer E/S0s are observed than predicted implying a net negative evolutionary process --- luminosity dimming, disassembly or masking by dust --- at $I > 24$. The breadth of the redshift distribution at faint magnitudes implies strong clustering or an extended epoch of formation commencing at $z > 3$. \\noindent (2) Spiral galaxies are present in numbers consistent with zero-evolution predictions to $I = 22$. Beyond this magnitude some net-positive evolution is required. Although the number-counts are consistent with the passive-evolution predictions to $I=26.0$ the redshift distributions favor number and luminosity evolution although few obvious mergers are seen (possibly classified as Irregulars). We note that beyond $z \\sim 2$ very few ordered spirals are seen suggesting a formation epoch of spiral galaxies at $z \\sim 1.5$ -- $2$. \\noindent (3) There is no obvious explanation for the late-type/irregular class and this category requires further subdivision. While a small fraction of the population lies at low redshift (i.e. true irregulars), the majority lie at redshifts, $1 < z < 3$. At $z > 1.5$ mergers are frequent and, taken in conjunction with the absence of normal spirals at $z > 2$, the logical inference is that they represent the progenitors of normal spirals forming via hierarchical merging. ", + "introduction": "The Hubble Deep Field (HDF; Williams {\\it et al.} 1996) has provided the deepest and clearest window to date on the extragalactic sky. From this dataset, groups have studied the morphologies of the faintest galaxies (e.g. Odewahn {\\it et al.} 1996; Abraham {\\it et al.} 1996) and made photometric estimates of the redshift of these objects (e.g. Lanzetta, Yahil \\& Fern\\'andez-Soto 1996; Brunner {\\it et al.} 1997). Here we combine these two independent analyses to generate morphological number-counts {\\it and} morphological redshift distributions for a complete sample of objects from the Hubble Deep Field (413 objects to $I=26.0$). This represents a unique dataset which provides strong constraints on the many faint galaxy models which have been postulated to explain the phenomena of the faint blue excess (see Ellis 1997 for a recent review). Faint galaxy models fall into three broad generic categories: dwarf-dominated models; pure luminosity evolution models; and merger models. All of these various models can provide a fit to the observed faint galaxy number-counts and therefore these data alone are insufficient to distinguish between the proposed models. Additional observational constraints are required and the most definitive one is that of the observed redshift distributions, N(z), for progressively fainter magnitude slices. For example, dwarf-dominated models (Driver {\\it et al.} 1994; Phillipps \\& Driver 1995; Babul \\& Ferguson 1996) predict an additional low redshift component at faint magnitudes when compared to the N(z) predictions of the zero-evolution models. Conversely pure-luminosity evolution models (e.g. Metcalfe {\\it et al.} 1995; Campos \\& Shanks 1997) predict a high redshift component whilst merger models lie somewhere in between (e.g. Carlberg 1992; Rocca-Volmerange \\& Guiderdoni 1990). In theory then, the problem is surmountable; in practice obtaining a comprehensive and complete spectroscopic redshift distribution at faint magnitudes is beyond our technological capabilities. The very faint redshift surveys which do exist (e.g. Glazebrook {\\it et al.} 1995a; Cowie {\\it et al.} 1996) are relatively small samples and arguably susceptible to selection biases (e.g. wavelength coverage, spectral features, surface brightness). For the moment the only recourse for establishing the N(z) distribution at these faint magnitudes is to utilize distance estimates based on multi-band photometry, {\\it i.e.} photometric redshifts. In \\S 2 we briefly discuss the adopted morphological and photometric HDF catalogs. In \\S 3, we compare the resulting galaxy number-count data and redshift distributions to zero- and passive- evolution models, and infer the generic form of evolution implied by the data. \\S 4 summarises our findings. ", + "conclusions": "We have combined the morphological catalog of Odewahn {\\it et al.} (1996) with the photometric redshift catalog of Lanzetta {\\it et al.} (1996) for all objects in the Hubble Deep Field to $I = 26$. This has resulted in a unique dataset from which we can construct the observed number-counts {\\it and} redshift distributions for E/S0s, Sabcs and Sd/Irrs down to $I = 26$. Adopting the local morphological luminosity functions (Marzke {\\it et al.} 1994) and with the caveat of {\\it a uniform overall renormalization at $b_{j} = 18$}, we conclude the following: Ellipticals form over an extended period starting at $z > 3$, however the observed underdensity in the number-counts implies that young ellipticals are either masked by dust or only become recognizable morphologically as ellipticals after their stellar population has stabilized and aged ({\\it i.e.} a substantial population of young overly luminous ellipticals is not seen). From the observed absence of $L_{*}$ spirals at moderate to high redshifts ($z > 2.0$) we conclude that present-day disks are forming at $z \\sim 2$ via hierarchical merging. During this stage their morphologies are highly irregular, this is corroborated by the high number of irregulars seen at this epoch. At lower z the merger rate sharply declines and the more luminous (massive ?) objects crystallize into the regular spiral systems and evolve passively with minimal further merger events. Meanwhile the remaining less luminous disk systems and merger by-products/remnants fade ($z > 1$) into the local dwarf and low surface brightness populations. Our final {\\it Hubble Deep Field} catalog of morphologies and photometric z's is available on request from: spd@edwin.phys.unsw.edu.au" + }, + "9802/astro-ph9802236_arXiv.txt": { + "abstract": "We present a new scenario for the development of the Universe after the Big Bang, built on the conjecture that a vast majority of the primordial quark matter did not undergo a phase transition to normal nuclear matter, but rather split up into massive quark objects that remained stable. Hence, such primordial quark matter would make up the so-called dark matter. We discuss, mostly in qualitative terms, the consequences for galaxy formation, the origin of normal matter, the occurrence of massive black-holes in galactic centres and the cosmic gamma-ray bursts. ", + "introduction": "One of the most fascinating mysteries in modern astrophysics and cosmology is the nature and origin of the so-called dark matter in the cosmos. It is (by definition) non-luminous, and reveals itself only through its gravitational interaction with the luminous galactic matter or with light. Its main signature is that most studied galaxies rotate in a ``non-Keplerian'' way, as estimated from their luminous matter \\cite{Raffelt97}. It appears as if the galaxies contain some extra, non-luminous matter, with a total mass believed to be about an order of magnitude higher than that of the luminous matter. There are weaker indications that also galaxy clusters behave strangely, and therefore would contain some extra matter inbetween the galaxies. There are also theoretical, cosmological arguments favouring an overall densitity much higher than the one estimated from direct observations of luminous matter. Detailed studies of galaxy rotations seem to indicate that dark matter has a more extended density profile than the luminous matter, stretching out to several times the conventional galactic radii. However, there is no general agreement as to the geometrical shape of the dark-matter halo of a typical spiral galaxy, and suggestions of spherical as well as slightly flattened mass distributions can be found in the literature. Neither is there a generally accepted functional dependence of the density $\\rho (r)$ as a function of the distance $r$ to the galactic centre, even for fits to spherical distributions. One often assumes a form that at least asymptotically falls off as $r^{-\\alpha}$. In most analyses an $\\alpha < 2$ is used, so that one has to introduce a cut-off in $r$ in order to avoid an infinite galactic mass; {\\it i.e.}, an {\\it ad hoc} galactic radius. Normally, one assumes that the total mass of a galaxy is around ten times that of the luminous matter, while the radius can vary widely, depending on the exact choice of density distribution. Some recent discussions of best-fit analytical forms can be found in \\cite{Moore97}. The widely varying forms originate from different analytic methods (except for the trivial reason that they are sometimes used for different galaxies). Different groups aim at fitting different parts of the dark-matter distribution, which sometimes refers to all galactic dark matter, and sometimes only to the part being more peripheral than our solar system. There is also a distinction between analyses that are founded on model-dependent simulations of galaxy formation, including conjectured values of cosmological parameters, and those built on observations of the Milky Way and other galaxies. During the last few years an extensive study of stars in the Large Magellanic Cloud (LMC) has also found a few cases of gravitational lenses, so-called machos (Massive Compact Halo Objects), which magnify the light from those stars, and seem to move in the outskirts of our galaxy. The masses of the discovered objects lie below the solar mass ($M_{\\odot}$), maybe in the range $(0.01-0.8) M_{\\odot}$ \\cite{Alcock96}. It is an open question if the machos can make up for all dark matter in the Milky Way. The MACHO collaboration itself gives some support to an average value for the macho mass of $0.5M_{\\odot}$, and a $50$ {\\it per cent} macho fraction in the Milky Way halo (assuming a spherical shape). There is no lack of imaginative models for the dark matter, a majority of which rely on ideas that have never been confirmed, or even supported, by independent experiments or measurements on earth or in space \\cite{Raffelt97}. The least speculative ones are those built on astrophysical ideas about dark, compact objects of normal matter, {\\it e.g.}, bodies created like normal stars, but with too small masses to ignite fusion processes and become luminous. So-called jupiters and brown dwarfs fall into this category. One can also think of planets or comets that have escaped solar systems in large numbers. Most explanations built on particle physics are much more speculative. Sometimes they even rely on new fundamental ideas that were invented {\\it just} to explain dark matter. Examples of interesting, but entirely hypothetical, particles assumed to contribute to the dark matter, are ``neutralinos'' and ``axions'' \\cite{Raffelt97}. There are only two particle-physics motivated dark-matter models that are based on relatively well-known, or theoretically well-studied phenomena. The simplest one is the ``heavy-neutrino'' model, {\\it i.e.}, the suggestion that at least one neutrino species has a rest mass high enough to make up for the galactic dark matter. Such neutrinos cannot, of course, explain the discovered gravitational lenses in the Milky Way, neither can they easily be reconciled with the fact that dark matter seems more peripheral than normal matter. There are also more general problems with galaxy formation, and it is believed that heavy neutrinos can account for only a small fraction of the dark matter \\cite{Sarkar97}. The other model of this kind identifies dark matter with objects consisting of a so-called {\\it quark-gluon plasma} (QGP), {\\it i.e.}, a form of matter with only quarks, and no structuring into protons and neutrons. Such objects are expected from, or at least not forbidden by, basic quark theory; quantum chromodynamics (QCD). There is no reason to believe that systems of just a few quarks (two or three) would be the only ones of physical or astrophysical interest. There is indeed an intense current research about QGPs. One example is the extensive experimental programmes at several high-energy laboratories, which aim at creating a QGP in heavy-ion collisions. The idea is to compress nuclear matter into such a dense state that individual nucleons can no longer be distinguished. Then the QGP might perhaps survive long enough to send out some clear signals, before converting (``hadronising'') into normal matter again. A few hints of QGP creation have indeed been found (``strangeness enhancement'' and ``$J/ \\Psi$ suppression''), above all at the CERN laboratory in Geneva, but these have also been disputed, and claimed to be consequences of more conventional physical phenomena. The extension of QGP ideas into astrophysics is straightforward, since the bulk of matter must have been in the form of such a plasma just after the Big Bang, when densities were still far above those inside atomic nuclei, and of the order of those inside protons and neutrons. There are also many current astrophysical situations where one can think of extremely high densities, such as inside a neutron star, or a collapsing would-be supernova. Normal hadronic matter is believed to have been created spontaneously as soon as the Universe expanded into an average density of the order of normal nuclear densities, although there is no agreement as for the details of this universal hadronisation. One can think of an explosive and very fast phase transition, going from the cooler periphery of the Universe and inwards, or a slower growth of bubbles of normal matter inside local density fluctuations, until finally the QGP disappeared, as shrinking bubbles inside the normal hadronic matter. It seems a very natural line of thinking to speculate that something went wrong within this scenario, so that a vast majority of the primordial matter stayed in the QGP phase, and now constitutes the dark matter. From a minimalistic point of view it is more natural to build on the fact that the Universe has {\\it already been} in a specific ``dark-matter'' state, than to speculate about completely unknown forms of matter. However, this requires the QGP to be {\\it the absolute ground state of matter}, at least in some cosmically interesting mass region. Ideas along such lines began to flourish in the early 1970s, with a pioneering work by Bodmer \\cite{Bodmer71}, and subsequent analyses in 1979 by Chin and Kerman \\cite{Chin79}, and Bjorken and McLerran \\cite{Bjorken79}. In 1984 De R\\'ujula and Glashow \\cite{Rujula84}, Fahri and Jaffe \\cite{Fahri84}, and Witten \\cite{Witten84} presented more refined analyses, which strengthened the conclusion that quark-matter ``nuggets'' or ```bags'' are indeed the ground state of matter, and hence contribute to dark matter. This idea cannot be rigorously underbuilt by basic principles, {\\it i.e.}, from QCD. Rather, the rule of the game is to rely on ``QCD-inspired'' phenomenological models, and the one most frequently used for analysing existing and hypothetical multi-quark objects is the so-called MIT bag model \\cite{Chodos74}. It is built on the assumption that quarks are confined to hadronic ``bags'' due to an external ``vacuum'' pressure, quantified by the so-called bag constant $B$, which takes a universal value, normally given as $B^{1/4} \\approx 150$ MeV. The $B$ value is fitted to known properties of normal hadrons, in a variational procedure, where the total energy (mass) of a hadron, primarily the proton, is minimised with the help of the bag radius. The known values of the proton mass and radius are then used to fit the model parameters, $B$ being the most crucial one. Many versions of the MIT bag model have been developed, containing various corrections, one of which is for interactions among the quarks, which were assumed to be free inside the bag in the original version. When analysing bags with more than three quarks, it appears as if already those with six quarks have a chance of being stable against decays via strong interactions. The case of such ``$H$ dibaryons'' \\cite{Jaffe77} is still debated in the literature, but no experimental evidence has so far been found. There is a clear trend within the model toward a higher stability for heavier objects with even more quarks. Such a stability appears in the analysis as a lower total energy per quark than inside a single nucleon. However, this conclusion requires that heavy quark objects, and already the $H$ dibaryon, contain an equal (or almost equal) number of the three lightest quarks, the $u$, the $d$ and the $s$ quarks. In this respect, the quark objects predicted by the MIT bag model differ from the structure of a hypothetical, strongly compressed atomic nucleus. The reason that the $s$ quarks are so crucial, in spite of their assumed higher mass, is that the Pauli principle does not forbid them to occupy the low-energy states of the $u$ and $d$ quarks. Once the $s$ quarks are captured in low-lying energy levels they cannot decay through weak interaction, much in the same way as neutrons in atomic nuclei can be stable. Still, some authors present results for a QGP of just $u$ and $d$ quarks, but this is mainly motivated with simplicity arguments in order to avoid computational difficulties. Hence, the macroscopic quark objects predicted to be of astrophysical value are often referred to as strangelets, strange stars, etc. There are several obvious, but interesting, consequences of this peculiar composition. One is that a quark object can be electrically neutral, due to the charges $(+2e/3,-e/3,-e/3)$ of the $(u,d,s)$ quarks. Hence a ``strangelet'' qualifies as dark matter, since it cannot attract an electron cloud, and therefore not send out light through ``atomic'' processes. It can still emit temperature radiation would it have a hot surface, but if it represents the ground state of matter, it will do so even at temperature $T = 0$. Another observation is that a strangelet cannot easily be produced by a contraction of atomic nuclei in high-energy heavy-ion reactions at accelerators. It would require weak-interaction conversions to $s$ quarks of many of the original $u$ and $d$ quarks in the very short time available before the nuclear matter flies apart again. At the best, one can hope for a ``baryon-free'' plasma, consisting of newly created quarks and antiquarks. If so, one might detect an enhanced production of hadrons containing $s$ and $\\bar{s}$ quarks, as the result of a limited QGP creation. The experimental evidence for enhanced strangeness due to such processes is under an intense study \\cite{Greiner98}. However, an almost exact balance between the three quark species is expected inside the cosmological plasma created after the Big Bang. In fact, ``any'' set of equally fundamental particles were in balance as long as the temperature (given in conventional mass units) was still much higher than the particle masses. Therefore, it suffices that the temperature obeyed $T \\leq 150$ MeV at the time of the $s$ quarks being captured inside their final quark objects. More massive quarks ($c$, $b$ and $t$) are not believed to exist in significant numbers inside absolutely stable quark matter, although they certainly contributed to the processes in the very early Universe. In the early work on such primordial quark nuggets in the cosmos, the authors did not commit themselves to a certain size (mass) of the objects, nor to an estimate of their absolute importance as dark matter. This probably has to do with the fact that numerical results from the MIT bag model, for objects containing more than a handful of quarks, are unreliable. The original results for protons and other low-mass hadrons were achieved through exact solutions of the Dirac equation inside a spherical bag. Various approximations must be applied for systems with dozens of quarks, not to mention the $10^{30}$ quarks, or more, inside astrophysical objects. Many of these are built on computational methods from nuclear physics, or from relativistic statistical mechanics. One typical MIT bag model analysis of this kind, presented in \\cite{Greiner98}, shows that a QGP with less than about $20$ quarks has a higher total energy than a corresponding nucleus, or set of free nucleons. This result assures that light nuclei do not decay to a QGP. However, heavier objects have lower total energies per quark, and are hence stable according to this model. This has been analysed up to a few hundred quarks, where the limit is set by practical computer capacities. We will assume that these MIT bag model results apply all the way up to quite heavy objects. For extremely heavy (astrophysical) objects also gravity is assumed to play an important role, and the MIT bag model must be complemented with general relativity. We will discuss this problem later. Strangely enough, this simple and logical dark-matter model does not seem to have reached a popularity in line with those built on much more exotic and controversial ideas. The modern literature on a quark-object dominated Universe is scarce, and, in fact, non-existent (would it not be for a persistent interest in its very early stage, and the presumed phase-transition from quark-matter to hadrons). This neglect cannot have been caused by published counter-arguments, because these are scarce too. A much quoted one is an analysis by Madsen \\cite{Madsen88}. He studied the effect of quark nuggets in space hitting a neutron star, and found that an impact of a very small amount of strange QGP into the dense interior of a normal neutron star would catalyse a phase transition of its full core. Hence, even a tiny lump of quark matter would turn a neutron star into a ``quark star'' (or ``hybrid star''), with a QGP-dominated interior. Such an extremely compact core would prevent the neutron star from experiencing a so-called pulsar glitch, {\\it i.e.}, a sudden change of rotational frequency, which is believed to be caused by a ``starquake'' coupled to an immediate change of the moment of inertia around the rotational axis. Since the probability of a pulsar glitch occurring in a neutron star can be estimated from observations, these arguments lead to an upper limit for the density of quark-matter nuggets in space, which, according to Madsen, excludes them as the dominant dark matter. However, Madsen's conclusion has been disputed, and there are even claims that pulsar glitches {\\it require} the presence of a quark-matter core, or are caused by the very phase transition when the core turns from neutron to quark matter. Hence, there is by now a rich literature built on the idea that some, or all, neutron stars have quark-matter cores. In fact, the main interest in cosmic quark matter in the modern literature now focuses on such conventional compact stars. Here, the quark matter is assumed to appear as the result of a gravitational contraction of normal matter, or of a rapid collapse after a supernova explosion. A review of the current literature can be found in Glendenning's book {\\it Compact Stars} \\cite{Glendenning96}. The true reason for quark matter to be out of fashion as a dark-matter candidate therefore seems to be sociological, namely that other ideas are more in line with the current development in theoretical and experimental high-energy physics. Here, the trends are toward ``smaller length-scales'' and higher mass-scales, parallel to the construction of new accelerators at CERN and elsewhere. This has led to a strongly increased interest in concepts like supersymmetry, leptoquarks, massive neutrinos and Higgs bosons, including a full spectrum of interesting astrophysical and cosmological implications. An important reason is probably also that some of these ideas, {\\it e.g.}, on massive neutrinos and on supersymmetric partners of quarks and leptons, can be tested with the impressive astrophysical Cherenkov detectors now in use around the globe. Dark-matter quark objects would be more elusive in this sense, primarily because they would be orders of magnitude less frequent than the exotic single-particle candidates, due to their much higher masses. Also, the atmosphere (or the water, ice or rock surrounding the Cherenkov detectors) would presumably erase their traces. Such traces are not even well-defined, and can therefore probably not be discriminated from an impact of normal atomic nuclei in, {\\it e.g.}, a space-born detector. Nevertheless, we argue that the original ideas of quark objects as dark matter are worthy of a revival, as they have not been convincingly counter-proved, and since they are built on a very simple principle, {\\it i.e.}, the one about the absolute stability of massive multi-quark states. We will assume that practically all dark matter is in this form, and try to limit the model parameters with the help of astrophysical data. For natural reasons, the discussion will be mainly qualitative since really conclusive data are indeed scarce. Also, the particle-physics foundations are not too well known, and in particular not the proper way of using the MIT bag model for very massive objects. ", + "conclusions": "The model presented here is, to the best of our knowledge, the first one that relates gamma-ray bursts to the dark-matter problem, and to the creation of normal atomic matter in the galactic centres. These ideas limit the model in the sense that they require a completely different sequence of events immediately after the Big Bang, as compared to the standard scenario. The most original ``details'' here are that the global quark-gluon plasma must have gathered into proto-galaxies {\\it before} the quark-hadron phase transition, and that the final structure of the galaxy can be a result of an explosive development, following the creation of a massive, central black hole, or a fireworks of early mergers of quark-object binaries. Such a fantastic scenario does not seem to violate qualitative astrophysical facts, but certainly most of the more conventional models, including the later phases of the so-called standard Big Bang scenario. It remains to analyse if it is also in line with the wealth of more detailed data from the cosmos. In the longer perspective it will also be interesting to wait for some clear-cut signatures of the more than $1000$ billion massive quark objects that we expect to orbit our own galaxy {\\it e.g.}, in the form of a much better statistics in the studies of gravitational micro-lenses. This project is supported by the European Commission under contract CHRX-CT94-0450, within the network \"The Fundamental Structure of Matter\"." + }, + "9802/astro-ph9802285_arXiv.txt": { + "abstract": "We discuss the implications of the recent discovery that Starburst Nucleus Galaxies (SBNGs) have lower oxygen abundances than ``normal galaxies\" of the same morphological type. Our interpretation of this result is that SBNGs are young galaxies, still in the process of formation. We consider several alternatives, but none of them provides a viable explanation. This new result has important consequences for our understanding of galaxy evolution, as it confirms the scenario of hierarchical formation of galaxies, and explains recent observations of the {\\it Hubble Space Telescope}. ", + "introduction": "The current paradigm about starbursts is that they are sporadic events that can occur at any time in the evolution of a galaxy, generally triggered by galaxy interactions (Heckman 1997). One of the reasons for this assumption is probably that starbursts are often seen in gravitationally interacting galaxies, and this paradigm has been reinforced in the last decade by the discovery that ultraluminous infrared galaxies are generally violently interacting. In a recent paper (Coziol et al.\\,1997a), we gathered published and unpublished data on several large samples of Starburst Nucleus Galaxies (SBNGs) to show that they have lower oxygen abundance than normal galaxies of the same morphological type, and that early-type SBNGs are even more deficient in oxygen than late-type ones. The simplest interpretation of this result is that SBNGs are young galaxies, which are still in the process of formation. The fact that early-type galaxies are less abundant in oxygen than late-type ones can be readily understood in the frame of the theory of hierarchical formation of galaxies. In this scenario of galaxy evolution (Tinsley \\& Larson 1979; Struck-Marcell 1981), ellipticals and bulges of spirals are formed by mergers of stellar and gaseous systems, while disks form by the collapse of the gas left from the successive mergers. ", + "conclusions": "" + }, + "9802/astro-ph9802291_arXiv.txt": { + "abstract": "Recent results have questioned the description of the QSO luminosity function in terms of a pure luminosity evolution and call for a luminosity dependent luminosity evolution. Measurements of the QSO clustering amplitude and evolution allow further distinguishing among the various physical scenarios proposed to interpret the QSO phenomenon. The general properties of the QSO population would arise naturally if quasars are short-lived events connected to a characteristic halo mass $\\sim 5 \\cdot 10^{12}$ M$_\\odot$. This is the typical mass of groups of galaxies in which the interactions triggering the QSO activity preferentially take place. ", + "introduction": "QSOs have played an important role as cosmic probes of the young universe: they are used as light beacons for absorption-line and gravitational lensing studies, as markers of galaxy formation activity and have guided astronomers in hunting up primeval galaxies. Their very existence at high redshifts is a challenge for cosmological models and historically the evolution of the QSO population has been one of the first evidences for an evolving universe. The nature of QSOs and the causes of their evolution, however, are far from being fully understood. In the standard ``demographic'' approach the basic physical ideas are tested against the observed shape and evolution of the QSO luminosity function (LF). As usual, with the growth and improvement of the databases this type of interpretation has become significantly more complex and we will argue that to disentangle the physical evolutionary patterns of QSOs additional information is needed, for example about their clustering properties. ", + "conclusions": "" + }, + "9802/astro-ph9802258_arXiv.txt": { + "abstract": "We analyze the spherically symmetric Einstein field equation with a massless complex scalar field. We can use the Newtonian solutions to fit the rotation curve data of spiral and dwarf galaxies. From the general relativistic solutions, we can derive high gravitational redshift values. ", + "introduction": " ", + "conclusions": "" + }, + "9802/astro-ph9802222_arXiv.txt": { + "abstract": "Gamma-ray bursts are sudden releases of energy that for a duration of a few seconds outshine even huge galaxies. 30 years after the first detection of a gamma-ray burst their origin remains a mystery. Here I first review the ``old'' problems which have baffled astronomers over decades, and then report on the ``new'' exciting discoveries of afterglow emission at longer wavelengths which have raised more new questions than answered old ones. ", + "introduction": "Gamma-ray bursts (GRBs) were first detected in 1967 with small gamma-ray detectors onboard the Vela satellites (Klebesadel \\etal\\ 1973) which were designed to verify the nuclear test ban treaty between the USA and the USSR. For many years the prevailing opinion was that magnetic neutron stars (NS) in the galactic disk were the sources of GRBs. No flaring emission outside the gamma-ray region could be detected, and no undisputable quiescent counterpart to a GRB could be established. Despite a distance ``uncertainty'' of 10 orders of magnitude, numerous theories (see a compilation in Nemiroff 1994) were advanced to explain the source of energy in GRBs. The measurements since 1991 of the Burst and Transient Source Experiment (BATSE) onboard the Compton Gamma-Ray Observatory have shown unequivocally that GRBs {\\bf are} isotropic even at the faintest intensities, and that there {\\bf is} a distinct lack of faint bursts as compared to a homogeneous distribution. An unprecedented wealth of additional information on each burst could be collected, yet the GRB origin remained a mystery. Over the previous two decades, GRB coordinates came with two mutually exclusive properties: {\\it arcmin accuracy} as provided by the interplanetary network (Hurley 1995) or {\\it fast} as provided by the BATSE Coordinate Distribution Network (BACODINE) system (Barthelmy \\etal\\ 1996). Only since the launch and successful operation of the Italian/Dutch BeppoSAX satellite is it possible to obtain accurate GRB positions in reasonably short time (few hours) which allow quick follow-up observations (Heise \\etal\\ 1998). The discovery of X-ray afterglow emission with the BeppoSAX satellite and related optical and radio transients has given a dramatic boost to both observations and theoretical investigations of GRBs over the last few months. At the present time (late 1997), our knowledge is evolving extremely rapidly. Thus, it may not be surprising that the content of this review has been expanded considerably as compared to the oral version given in May 1997. As in the talk, I will not cover Soft Gamma Repeaters, reviews of which can be found in Kulkarni (1998) and Smith (1998). ", + "conclusions": "The discovery of long-wavelength afterglow emission lasting for days to months constitutes a turning point in GRB research. However, despite these exciting new data most of the basic questions remain unanswered. While a cosmological distance scale seems to be generally accepted, the nature of the GRB host remains open. Furthermore, the complexity of GRB time histories, duration and spectra as well as spectral evolution during the bursts are not easily explained in the various variants of the fireball model which is widely accepted as the main scenario for the production of the $\\gamma$-ray burst emission. Finally, the origin of the GRB energy source remains a mystery. Possibly, knowledge of the history of the early cosmic evolution is required to understand the origin of GRBs. Thus, despite (or due to) the discovery of flaring counterparts at longer wavelengths GRBs remain an exciting field of research in the foreseeable future." + }, + "9802/astro-ph9802364_arXiv.txt": { + "abstract": "Effective star formation rates in tabular form are computed which yield a prescription for the star formation activity in model galaxies as a function of ambient density, metallicity, and stellar feedback. The effects of supernova explosions on the thermal balance of the Interstellar Medium (ISM) and the presence of a multi-phase ISM are explicitly included. The resulting grid of models can be implemented easily in N-body codes for the computation of star formation processes in merging galaxies and cosmological simulations. ", + "introduction": "The past decade has seen tremendous advances in the field of numerical cosmology (Navarro, Frenk \\& White 1996 and references therein). N-body simulations and hydro codes have benefitted from the large increase in CPU and memory capabilities of the current generation of supercomputers, and have reached a high enough sophistication to describe the development of large scale structure in the universe and the formation of galaxies with redshift. An important issue which ultimately needs to be addressed, in order to compare theory with observations of high redshift galaxies, is the formation of stars in proto-galactic structures (Norman \\& Spaans 1997, hereafter NS97, and references therein). Many parameterizations of the star formation rate in the ambient Interstellar Medium (ISM) exist in the literature and mostly derive from a Schmidt law applied to a sufficiently large body of gas (Theis et al.~1992; Kauffmann \\& White 1993; Spaans \\& Norman 1997, hereafter SN97). Star formation is a local phenomenon which must find its explanation in the stability and fragmentation of dense molecular clouds. Studies in our own Galaxy have focussed on the structure of dense proto-stellar cores, along with the chemical and thermal balance of star-forming regions (Helmich 1996, and reference therein). These studies lend indirect support to a Schmidt law, but emphasize the need to include explicitly the structure of the multi-phase ISM to model accurately the most important heating and cooling processes. A large unknown in these investigations is the relative importance of feedback. Supernova explosions and stellar radiation associated with the process of star formation influence the global physical structure of the interstellar gas which supports this process. Ideally, in the cosmological context one would like to solve for the properties of the thermal and chemical balance of the ISM and the star formation rate simultaneously with the solution of the gravitational N-body and hydrodynamical problem. Such an approach is beyond what is currently feasible and one is forced to decouple the ISM and star formation problem from the dynamical one. Along these lines, the purpose of this work is to present a set of numerical simulations which span a large range in density and metallicity, include a quantitative prescription for feedback, and yield the resulting star formation rate in the ambient medium in tabular form. This work extends the methods developed in a series of papers where the structure of the ISM and the importance of feedback effects have been investigated for proto-galactic disks (NS97), dwarf galaxies (SN97), massive proto-spheroidals in the Hubble Deep Field (Spaans \\& Carollo 1997), and elliptical galaxies (Carollo \\& Spaans 1998). The strength of the procedure used in this work to produce the effective star formation rates is its detailed treatment of molecular line cooling, chemistry, and thermal balance (section 2) for a wide range of physical conditions. This allows for a detailed treatment of the multi-phase structure of the ambient ISM, and the feedback of stellar photons and supernova ejecta on its physical balance. The star formation rate is computed in our approach through a Schmidt law applied to the molecular phase of the ISM. The use of this empirical law presumes that star formation as observed in our local environment is representative of star formation in general. This seems a reasonable approach for spiral galaxies (Kennicutt 1989). In the following sections, the input physics are discussed. The processes involved are complex, and various physical phenomena influence the final outcome of the computation. Therefore, one should keep the following major calculational steps in mind. 1) The thermal and chemical balance of the medium is determined locally and used to compute the amount of molecular gas. 2) This allows a calculation of the formation rate of stars according to a Schmidt law and an initial mass function (IMF). 3) The produced radiation is propagated across the computational grid in a radiative transfer calculation, and used to determine heating, ionization and dissociation rates. 4) The supernova blast waves, produced as the end products of stellar evolution, are propagated across the grid as well and yield the local input of kinetic (gas bulk motions) and thermal (gas heating) energy, as well as metals. 5) Steps 3 and 4 provide input for step 1. This entire procedure is followed for many different ambient densities, to derive the enrichment and physical state of the gas together with the corresponding star formation rate. ", + "conclusions": "We have derived star formation rates in tabular form as functions of the ambient density, metallicity, and stellar feedback in proto-galaxies with a given total gas mass. We have explicitely included the effects of supernova explosions on the thermal balance of the ISM, and the transition to a multi-phase ISM. The strength of the method used to produce the effective star formation rates is the detailed treatment of molecular line cooling, chemistry, and thermal balance. The tabulated star formation rates can be implemented in N-body codes for the computation of star formation processes in merging galaxies and cosmological simulations. The accuracy of the presented numbers is expected to be of the order of a factor of two, and in any case the rates should not change by more than a factor of four. The full set of numerical results is included in Table 1, and can be obtained in machine readable form from MS. Although we have included the effects of feedback self-consistently, the present calculations cannot capture the effects of massive shocks associated with a large collection of SN explosions propagating through the ISM. A complete treatment of the latter requires the incorporation of the numerical scheme adopted here in a full hydro treatment. Investigations along these lines are being pursued." + }, + "9802/astro-ph9802152_arXiv.txt": { + "abstract": "We report the new finding that type 1 Seyfert nuclei (S1s) have excess [FeVII]$\\lambda$6087 emission with respect to type 2s (S2s). The S1s exhibit broad emission lines which are attributed to ionized gas within 1 pc of the black hole, whereas the S2s do not show such broad lines. The current unified model of active galactic nuclei explains this difference as that the central 1 pc region in the S2s is hidden from the line of sight by a dusty torus if we observe it from a nearly edge-on view toward the torus. Therefore, our finding implies that the coronal line region (CLR) traced by the [FeVII]$\\lambda$6087 emission resides in the inner wall of such dusty tori. On the other hand, the frequency of occurrence of the CLR in the optical spectra is nearly the same between the S1s and the S2s. Moreover, some Seyfert nuclei exhibit a very extended ($\\sim$ 1 kpc) CLR. All these observational results can be unified if we introduce a three-component model for the CLR; 1) the inner wall of the dusty torus, 2) the clumpy ionized region associated with the narrow line region at distance from $\\sim$ 10 to $\\sim$ 100 pc, and 3) the extended ionized region at distance $\\sim$ 1 kpc. ", + "introduction": "Optical spectra of active galactic nuclei (AGN) show often very high ionization emission lines such as [FeVII], [FeX], and [FeXVI] (the so-called coronal lines). Since the ionization potentials of these lines are higher than 100 eV, much attention has been paid to the understanding of the coronal line region (CLR) (e.g., Oke \\& Sargent 1968; Osterbrock 1977, 1985; Grandi 1978; Pelat et al. 1981; Penston et al. 1984). It is often considered that the CLR has an intermediate nature between the broad line region (BLR) and the narrow line region (NLR) because the high-ionization lines have critical densities for collisional excitation are of the order of $10^7$ cm$^{-3}$ and some Seyfert nuclei show CLR emission lines with FWHM $\\sim$ 1000 - 2000 km s$^{-1}$ (De Robertis \\& Osterbrock 1984, 1986; Veilleux 1988; Appenzellar \\& \\\"Osreicher 1988; Appenzellar \\& Wagner 1991). Recent photoionization model calculations have suggested that the CLR is located mostly within inner 10 pc (Ferguson, Korista, \\& Ferland 1997). In fact, Oliva et al. (1994) found a compact ($<$ 10 pc) CLR in nearby Seyfert, the Circinus galaxy, using the near infrared coronal line [SiVI] at 1.92 $\\mu$m. However, it is also known that some Seyfert nuclei have an extended CLR whose size amounts up to $\\sim$ 1 kpc [Golev et al. 1994; Murayama, Taniguchi, \\& Iwasawa 1998 (hereafter MTI98)]. The presence of such extended CLRs is usually explained as the result of very low-density conditions in the interstellar medium ($n_{\\rm H} \\sim 1$ cm$^{-3}$) makes it possible to achieve higher ionization conditions (Korista \\& Ferland 1989). The above complicated situation raises the question; {\\it Where is the CLR in AGN ?} According to the current unified models(Antonucci \\& Miller 1985; Antonucci 1993), it is generally believed that a dusty torus surrounds both the central engine and the BLR. Since the inner wall of the torus is exposed to intense radiation from the central engine, it is naturally expected that the wall can be one of the important sites for the CLR (Pier \\& Voit 1995). Recently, Gallimore et al. (1997) discovered a very compact ($\\sim$ 1 pc) ionized region in the S2 galaxy NGC 1068 in radio continuum. Since the inner radius of the accreting molecular ring traced by water vapor maser emission is $\\sim$ 0.5 pc (Greenhill et al. 1996), this ionized region seems indeed to be the inner wall of the torus. If the inner wall is an important site of CLR, it should be expected that the S1s would tend to have more intense CLR emission because the inner wall would be obscured by the torus itself in S2s. In order to examine whether or not the S1s tend to have the excess CLR emission, we study the frequency distributions of the [FeVII]$\\lambda$6087/[OIII]$\\lambda$5007 intensity ratio between S1s and S2s. ", + "conclusions": "The arguments described in the previous section suggest strongly that there are three kinds of CLR; 1) the torus CLR ($r < 1$ pc), 2) the CLR associated with the NLR ($10 < r < 100$ pc), and 3) the very extended CLR ($r \\sim$ 1 kpc). Therefore, if we take these three emission-line regions into account, we may have a unified picture for the CLR of AGN. Their basic physical properties are summarized in Table 1. A schematic illustration of the CLR is shown in Figure 3. We mention that there is the large scatter in the [FeVII]/[OIII] intensity ratio in both the S1s and the S2s. This scatter suggests that the contribution of CLR emission from the inner, extended, and very extended CLR may be different from object to object. Moreover, it should be remembered that a half of Seyfert nuclei show no evidence for the CLR (Osterbrock 1977; Koski 1978). This may be attributed to a gas-rich condition in the circumnuclear region, resulting in a lower ionization condition. This diversity may make it difficult to construct a simple photoionization model for the CLR as well as for the NLR itself (Ferland \\& Osterbrock 1986). In view of recent new observations and insights, we discuss the nature of the three kinds of CLR in AGN. 1) The torus CLR: Given the current unified model, it is naturally considered that clouds on the inner edges of dusty tori provide important sites for the CLR (Pier \\& Voit 1995). A typical electron density is estimated to be $10^{7-8}$ cm$^{-3}$ (Pelat et al 1981; Ferguson et al. 1997; Pier \\& Voit 1995). Since the emissivity of coronal lines is proportional to $n_{\\rm e}^2$ under conditions of $n_{\\rm e} < n_{\\rm cr}$, the torus CLR can be the most luminous component. This is indeed shown in Figure 1. We also have to note that iron is often depleted in the interstellar medium. However, since the inner edges of the tori are exposed to the intense radiation field from the central engine and thus dust grains may be destroyed (Pier \\& Voit 1995). This is another reason why the torus CLR is more luminous than that in the NLR. We also mention that the torus CLR consists of many small ionized gas clumps though we illustrate it as shown in Figure 3. If we assume that the inner wall obeys a Keplerian rotation, we obtain a typical line width FWHM $\\simeq 2 v_{\\rm rot} \\simeq 1320 M_8^{1/2} r_{1}^{-1/2}$ km s$^{-1}$ where $M_8$ is the dynamical mass within a radius $r_1$ in units of $10^8 M_\\odot$ and $r_1$ is the radius of the NLR in units of 1 pc (Pier \\& Voit 1995). This fiducial value is almost comparable with those of coronal lines whenever they are broad (De Robertis \\& Osterbrock 1984, 1986; Appenzellar \\& \\\"Osreicher 1988; Appenzellar \\& Wagner 1991; Giannuzzo et al. 1995). It has been known that some Seyfert nuclei and quasars have ionized regions whose line widths are a few 1000 km s$^{-1}$ (Brotherton et al. 1994; Mason et al. 1996). Since these line widths are intermediate between those of the NLR and the BLR, it has been suspected that there is an intermediate line region (ILR) in addition to the traditional NLR and BLR. We propose that this ILR is just located at the inner wall of the tori and thus the torus CLR may be associated spatially with the ILRs. 2) The CLR associated with the NLR (the clumpy CLR): The recent {\\it Hubble Space Telescope} observations of the NLR in a number of nearby Seyfert nuclei have shown that the NLR consists of a large number of gas clumps and thus the structure of the NLR turns out to be much more complex than what we thought (Wilson et al. 1993; Bower et al. 1995; Capetti et al. 1996a, 1996b). Therefore, we refer this CLR as the clumpy CLR. It is naturally expected that the cloud surface facing to the continuum radiation may be the major site of CLR. MTI98 found that there is no correlation between [FeVII] and an optical FeII feature at $\\lambda$4570 which is presumed to arise from warm neutral, or partially-ionized regions of gas clouds. Further, there is no correlation between [FeVII] and [OI] (Murayama and Taniguchi, in preparation). These properties imply that highly ionized gas clumps are decoupled from low-ionization gas clumps. We may interpret that the CLR arises from matter-bounded ionized clumps while low-ionization lines arises mostly from ionization-bounded gas clumps\\footnote{ Recently, measuring the electron temperatures of both O$^{2+}$ and N$^{+}$ regions in the extended emission-line regions of several Seyfert galaxies, Wilson et al. (1997) proposed that the [OIII] emission arises mostly from the matter-bounded clouds while the [NII] emission arises from ionization-bounded clouds because the observed electron temperature difference between the [OIII] and [NII] regions is too large (e.g., $\\sim$ 5000 K) to be accounted for in terms of photoionization of ionization-bounded clouds. On the analogy of this finding, we consider that the NLR also consists of the two kinds of clumps.}. Since the [FeVII] emitting region should be exposed to the harder and stronger radiation field than the [OIII] region, the radiation pressure exerted from the central engine is higher for high-ionization gas clumps than for the low-ionization ones and thus the high-ionization clumps may be accelerated more efficiently, leading to the systematic blueshift of the high-ionization clumps with respect to the lower ones. This property has been often observed in many AGN (Grandi 1978; Appenzellar \\& \\\"Osreicher 1988; Gaskell 1982; Wilkes 1984). 3) The extended CLR: The very extended CLR ($r \\sim$ 1 kpc) is found in both NGC 3516 (Golev et al. 1994) and Tololo 0109$-$383 (MTI980. If the interstellar medium consists of very low-density gas clouds (e.g., $n_{\\rm H} \\sim 1$ cm$^{-3}$), the high ionization condition can be achieved (Korista \\& Ferland 1989). Therefore, it is reasonable that some Seyfert nuclei have such the very extended CLR. The extended CLR seems to be related to the so-called extended ($r \\sim$ 1 - 10 kpc) narrow line region (ENLR; Unger et al. 1987) which is thought to be interstellar medium photoionized by the continuum radiation from the central engine. The CLR may be lower-density parts of the ENLR because the CLR needs higher ionization condition than the typical [OIII] emitting region. \\vspace {0.5cm} We would like to thank B. Vila-Vilalo for useful discussion. TM was supported by the Grant-in-Aid for JSPS Fellows by the Ministry of Education, Science, Sports, and Culture. This work was financially supported in part by Grant-in-Aids for the Scientific Research (No.\\ 0704405) of the Japanese Ministry of Education, Science, Sports, and Culture. \\clearpage \\begin{deluxetable}{lcccc} \\tablecaption{The three-component model for the CLR} \\tablehead{ & \\colhead{$r$ (pc)} & \\colhead{$n_{\\rm H}$ (cm$^{-3}$)} & \\colhead{FWHM (km s$^{-1}$)} & \\colhead{Associated} \\nl & & & & \\colhead{Emission-line region} \\nl } \\startdata Torus CLR & $\\sim$ 1 & $\\sim 10^7$ - $10^8$ & 1300\\tablenotemark{a} & ILR\\tablenotemark{b} \\nl Clumpy CLR & $\\sim$ 1 - 100 & $\\sim 10^3$ - $10^6$ & 400 - 750 & NLR\\tablenotemark{c} \\nl Extended CLR & $\\sim$ 1000 & $\\sim 1$ & $<$ 50\\tablenotemark{d} & ENLR\\tablenotemark{e} \\nl \\enddata \\tablenotetext{a}{An observed FWHM is 2$v_{\\rm rot}$sin$i$ where $i$ is the angle between the line of sight and the rotational axis of the torus. When we observe the torus from a face-on view, the FWHM should be a virial line width, 660 $M_8^{1/2} r_1^{-1/2}$ km s$^{-1}$. Such narrow line widths of the CLR are observed in some Seyfert nuclei (De Robertis \\& Osterbrock 1984, 1986; Giannuzzo et al. 1995).} \\tablenotetext{b}{Intermediate line region.} \\tablenotetext{c}{Narrow line region.} \\tablenotetext{d}{Unger et al. (1987)} \\tablenotetext{e}{Extended narrow line regions. Another term, extended emission-line regions (EELRs) is also used.} \\end{deluxetable} \\clearpage \\clearpage" + }, + "9802/astro-ph9802247_arXiv.txt": { + "abstract": "Magnitude--limited samples of spiral galaxies drawn from the Ursa Major and Pisces clusters are used to determine their extinction properties as a function of inclination. Imaging photometry is available for 87 spirals in $B,R,I$, and \\k' bands. Extinction causes systematic scatter in color--magnitude plots. A strong luminosity dependence is found. Relative edge-on to face-on extinction of up to $1.7^m$ is found at $B$ for the most luminous galaxies but is unmeasurably small for faint galaxies. At $R$ the differential absorption with inclination reaches $1.3^m$, at $I$ it reaches $1.0^m$, and at \\k' the differential absorption can in the extreme be as great as $0.3^m$. The luminosity dependence of reddening can be translated into a dependence on rotation rate which is a distance-independent observable. Hence, corrections can be made that are useful for distance measurements. The strong dependence of the corrections on luminosity act to steepen luminosity-linewidth correlations. The effect is greatest toward the blue, with the consequence that luminosity--linewidth slope dependencies are now only weakly a function of color. ", + "introduction": "It is evident that there is optical obscuration in spiral galaxies yet the quantitative amount of the dimming due to obscuration is disputed. There is a displacement in the photometric properties of edge-on and face-on galaxies. If distance considerations are taken into account, then in a set of similarly sized galaxies the more edge-on ones will tend to be fainter, or in a set of galaxies of the same luminosity the more edge-on ones will tend to be bigger. This opening statement touches on the two outstanding problems that confront efforts to quantify the level of obscuration in galaxies. To begin with, there is more than one good explanation for the photometric separation of galaxies with inclination: galaxies of the same intrinsic size and luminosity might appear dimmer toward edge-on because of increased path lengths through the obscuring material, or they might appear larger toward edge-on because surface brightnesses are enhanced by the increased path lengths. The second big problem arises out of uncertain relative distances with most samples since luminosities and dimensions scale differently with distance, hence distance uncertainties add noise to tests. A nice review of previous research is provided by Huizinga (1994)\\markcite{h94}. At one extreme, galaxies were accepted to be completely transparent and corrections were made only to dimensions in the {\\it Reference Catalogue of Bright Galaxies} (de Vaucouleurs \\& de Vaucouleurs 1964)\\markcite{rc1}. At the other extreme, galaxies were proposed to be optically thick at all observable radii by Valentijn (1990)\\markcite{v90}, to the extent that Gonz\\'alez-Serrano \\& Valentijn (1991)\\markcite{gv91} posited the obscuring material could resolve the `dark matter' problem. The current generally accepted viewpoint is intermediate. Spiral galaxies probably have high opacities near their centers but are relatively transparent at their visible extremities. We have a utilitarian need to understand the global obscuration properties of spiral galaxies. A primary consideration for us is the need to correct for inclination effects with luminosity -- HI profile linewidth distance estimators (Tully \\& Fisher 1977\\markcite{tf}). Inclination dependencies also have to be understood when we use photometric information in statistical studies of galactic structure (cf, Tully \\& Verheijen 1997\\markcite{tv}). In these cases, it may be enough to know the overall and statistical effects of projection from an empirical evaluation. An alternative approach is to try to develop a physical model of what is happening. However, a realistic model would have a lot of parameters and require a lot of information as constraints. There have been attempts in this direction (Wainscoat, Freeman, \\& Hyland 1989\\markcite{wfh}; Byun, Freeman, \\& Kylafis 1994\\markcite{bfk}; Bianchi, Ferrara, \\& Giovanardi 1996\\markcite{bfg}). To date, there has not been any study with both the information content per object needed for elaborate modeling and the large number of objects needed for a statistical evaluation of the norm and range of obscuration properties. What we bring that is new is imaging photometry over a substantial optical--infrared baseline: large format $B,R,I$ CCD imaging and \\k' imaging with a 256x256 HgCdTe detector. Our sample is not really large but it has nice qualities with respect to the distance problem and with respect to completion. We are using data from two clusters, hence objects are at either of two discrete distances, and there is a measure of completion to well specified limits. Our primary sample is drawn from the Ursa Major Cluster. The photometric data has been published by Tully et al. (1996)\\markcite{t5}. There are 62 galaxies in a window on the sky and in redshift that are brighter than $14.7^m$ at $B$. Optical photometry is available for all 62 and \\k' photometry is available for 60. Observations and reductions of the optical components were done by MJP, MAWV, and RBT. Observations and reductions of the \\k' material were carried out by JSH, MAWV, and RBT. The Ursa Major Cluster is a loose, irregular cluster with normal spirals and few gas-poor systems. It is probably dynamically young. The constituents may be representative of typical galaxies outside of dense regions. Our sample is complete down to roughly the luminosity of the Small Magellanic Cloud at $M_B \\simeq -16.5^m$. Unfortunately, there is a paucity of extremely luminous galaxies in this sample. Our supplemental sample is drawn from what has been called the Pisces Cluster, part of the Perseus-Pisces filament (Haynes \\& Giovanelli 1988\\markcite{hg88}), and an entity that has frequently been used in studies of distances and flows (Aaronson et al. 1986\\markcite{a86}; Han \\& Mould 1992\\markcite{hm}). Sakai, Giovanelli, \\& Wegner (1994)\\markcite{sgw} have discussed an overlapping part of the region and refer to subunits as the NGC~383 and NGC~507 groups. We consider candidates within the window $00^h 49^m < \\alpha < 01^h 32^m$, $28\\deg < \\delta < 34\\deg$, and $4,300~\\kms < V_{helio}+300{\\rm sin}\\ell{\\rm cos}b < 6,400~\\kms$. There are a couple of knots in the region dominated by early-type galaxies but spirals are scattered about in an irregular fashion. We take galaxies detected in the HI 21 cm line at Arecibo (Giovanelli \\& Haynes 1985\\markcite{gh85}, 1989\\markcite{gh89}, Wegner, Haynes, \\& Giovanelli 1993\\markcite{whg}), typed as spirals, and with cataloged axial ratios $b/a \\le 0.5$. The Arecibo Telescope gives a high detection rate for spirals at 5,000~\\kms\\ so there is reasonable completion (at $\\delta < 33.5\\deg$) to a magnitude limit of $B = 15.7^m$. This limit is $M_B = -19^m$ after corrections, roughly the magnitude of the Local Group galaxy M33. Cases where the HI signals may be confused are avoided. This Pisces sample has the disadvantages that our axial ratio limit excludes very face-on examples and for observational reasons it is not as rigorously complete as the Ursa Major sample (details in next section). However, it contains a substantial number of intrinsically luminous spirals. The optical photometry was acquired and reduced by MJP and RBT. The \\k' photometry was obtained by PLW, WS, and RBT. There are 38 galaxies in the Pisces sample. The information we have available allows us to excape the problems that have plagued efforts to calibrate the levels of obscuration in galaxies. We will describe tests that decouple luminosities and dimensions. The cluster nature of the samples essentially eliminates scatter due to uncertain relative distances. The sample completeness features are important. Two different tests have been considered. It turns out that the most sensitive test involves deviations as a function of inclination from the mean correlation in color-magnitude plots where one passband is $B$, $R$, or $I$ and the other passband is \\k'. Another test gives consistent but less statistically significant results. It involves deviations with inclination from the mean correlations between luminosities and HI profile linewidths. The latter test requires kinematic information as well as photometric and cannot include very face-on cases because of the large uncertainties in the deprojected kinematic parameter. Hence, fewer galaxies are available for the second test and the inclination baseline is truncated. Also, the ratio of reddening scatter to the intrinsic luminosity scatter in the relations is somewhat more favorable in the first test compared with the second. Between the two clusters, we have 101 galaxies with $B,R,I,$\\k' information in our magnitude limited sample. It will be seen that the statistical effects of obscuration are quite evident. The question becomes, how complex a description of the obscuration does the data afford? The simplest case would be the definition of a single coefficient for some analytic dependency on inclination or ellipticity in each bandpass. However, it has been suggested that there are dependencies on such additional parameters as type (de Vaucouleurs et al. 1991\\markcite{rc3}; Han 1992\\markcite{h92}) or luminosity (Giovanelli et al. 1995). Indeed, it will be shown that there is a strong dependence of obscuration on luminosity in our sample. As Giovanelli et al. (1995) point out, the correction procedure that is adopted affects the slope of the luminosity--linewidth correlation hence, potentially, distance estimates based on that correlation. The challenge to us is to generate a description of the luminosity dependence with our relatively limited sample. One needs enough cases in a bin to sample a broad inclination range and characterize the scatter. Then the luminosity dependence has to be transformed to a distance-independent parameter, like the profile width, to be useful for distance measurements. It would also be nice to further our understanding of how the obscuring material is distributed in galaxies. For example, Giovanelli et al. (1994)\\markcite{g94} and Peletier et al. (1995)\\markcite{p95} contend that there is a strong radial dependence, with considerable obscuration near the centers of galaxies and little at the outer edges. It would be possible to investigate that kind of effect with two-dimensional color decompositions with our data. However, that study is beyond the scope of this paper. ", + "conclusions": "Though our sample is moderate in number (pared to 87 spirals in the final analysis) it is reasonably well defined in terms of completion limits and we have photometric images in four bands ranging from $B$, where there is considerable absorption, to \\k', where there is minuscule absorption. Inclination effects of the sort attributable to extinction are seen in the scatter of linewidth--magnitude and color--magnitude correlations. More edge-on galaxies are fainter and redder in the mean. Plots of deviations from the mean as a function of inclination provide a quantitative measure of the amplitude of absorption. Similar results are found whether the residuals in linewidth--magnitude or color--magnitude relations are considered. The statistical significance of results are better with the color--magnitude analysis so we concentrate on those tests. The claim by Giovanelli et al. (1995)\\markcite{g95} is confirmed that the extinction in spirals is luminosity dependent. Galaxies with $M_{K^{\\prime}}^{b,k,i}$ reaching $-25^m$ ($M_B^{b,k,i}\\sim-22^m$) have up to $\\sim 1.7^m$ of differential absorption between face-on and edge-on ($1.3^m$ at $R$, $1.0^m$ at $I$, $0.25^m$ inferred at \\k'). However the amplitude of extinction drops off rapidly with decreasing luminosity, becoming unmeasurable at $M_{K^{\\prime}}^{b,k,i}>-20^m$ ($M_B^{b,k,i}\\ga-17^m$). The corrections reach higher values than we entertained in the past (Tully \\& Fouqu\\'e 1985)\\markcite{tf85}. We arrived at our old corrections from data on relatively nearby galaxies which have low luminosities in the mean. Samples such as those studied by Giovanelli et al. (1994,1995) have larger mean redshifts, hence larger mean luminosities, and it can now be understood why such samples lead to larger absorption corrections. Extinction corrections have been formulated in terms of the distance independent variable, linewidths, so they can be applied to distance estimators. The data assembled in this paper suggests $H_0=80$ \\kms Mpc$^{-1}$. \\bigskip Sadly, Peter Witchalls passed away a few days after this paper was submitted for publication. \\clearpage" + }, + "9802/astro-ph9802071_arXiv.txt": { + "abstract": "We have used archive HST WFPC2 data for three elliptical galaxies (NGC 3379 in the Leo I group, and NGC 4472 and NGC 4406 in the Virgo cluster) to determine their distances using the Surface Brightness Fluctuation (SBF) method as described by Tonry and Schneider (1988). A comparison of the HST results with the SBF distance moduli of Ciardullo et al (1993) shows significant disagreement and suggests that the r.m.s. error on these ground-based distance moduli is actually as large as $\\pm$0.25 mag. The agreement is only slightly improved when we compare our results with the HST and ground-based SBF distances from Ajhar \\etal (1997) and Tonry \\etal (1997); the comparison suggests that a lower limit on the error of the HST SBF distance moduli is $\\pm$0.17 mag. Overall, these results suggest that previously quoted measurement errors may underestimate the true error in SBF distance moduli by at least a factor of 2-3. ", + "introduction": "There now exist several methods of distance determination that appear to give measurements with small errors, \\ie the SBF method and the use of planetary nebula luminosity function (PNLF). Both of theses methods claim (and appear to deliver) distances with an accuracy of better than 10\\%. These methods are crucial in determining the zero-point and reliability of other distance estimators and also in determining the structure of the Universe in our neighbourhood out to about 30Mpc. In this paper we present an analysis of archive HST WFPC2 data of three ellipticals using the SBF method. The SBF method is based on the pixel-to-pixel brightness variations that are present in all galaxies, which is due to the varying number of stars within each pixel. The variations are proportional to $\\overline{f}\\sqrt{N_*}$ where $\\overline{f}$ is the mean flux of a star and $N_*$ is the average number of stars per pixel. To detect the SBF signal, it must be strong enough to be separated from sources of noise within the frame. The main sources of noise are shot noise and contamination from faint objects within the frame. The shot noise can be overcome by taking a long enough exposure, but the contamination from point sources needs better resolution to reduce its effect. Since the resolution of the HST is far superior to anything yet possible from the ground then images taken with the HST should reduce the point source contamination considerably. The HST data therefore will provide a crucial check of the accuracy of the ground-based SBF technique. \\begin{table} \\begin{tabular}{|c|c|c|c|c|} \\hline Galaxy & P.I. & Proposal ID & F814W & F555W \\\\ \\hline NGC 3379 & Faber & 5512 & 3x500,1x160 & 3x400,1x140 \\\\ NGC 4472 & Westphal & 5236 & 2x900 & 2x900 \\\\ NGC 4406 & Faber & 5512 & 3x500 & 3x500 \\\\ \\hline \\end{tabular} \\caption{HST archive data used for the SBF analysis. Exposure times for each filter are given in seconds.} \\label{ta:archive} \\end{table} \\section {The Observations} A search of the HST archive provided us with images of three galaxies, on which to perform the SBF analysis, with previous ground-based SBF distances which can then be checked. WFPC2 frames taken using the F814W and F555W filters were available for each of the galaxies NGC 3379, NGC 4472 and NGC 4406. The details of the frames are given in Table~\\ref{ta:archive}. All of the archive frames have the nucleus of the galaxy centred on the PC chip, with the outer regions present in all of the WF chips. On none of the frames were the central regions of the galaxy saturated. \\section {Data reduction} \\label{sec:reduc} All of the data were reduced using software available under IRAF\\footnote{IRAF is distributed by the National Optical Astronomical Observatories} and some specially written software by the authors. The frames for each galaxy and filter were co-added and any cosmic rays removed using the standard tools available under IRAF in the STSDAS package. \\subsection{Removing the galaxy background} Before we are able to investigate the SBF, the galaxy background needs to be removed from the images. For the PC images an ellipse fitting routine was used ({\\it ellipse} under the {\\it stsdas.analysis.isophote} package). This routine iteratively fits elliptical isophotes over a 2-dimensional image and produces an output which can be used to construct a smooth image of the fit. This image was then subtracted from the original image to give a resultant image which is flat on large scales. Generally the fitting worked well, but very close into the nucleus of the galaxy it failed. For some PC images the fitting broke down completely at the edges of the chip. For the WF images it was not possible to use any of the ellipse fitting routines available under IRAF since they all required that the center of the ellipse was in the image. Instead, we followed Simard \\& Pritchet (1994) and used {\\it fit1d} to provide a fit to the galaxy background. For each frame thirty {\\it fit1d} passes were made with five 2.5$\\sigma$ rejection iterations for each pass. For all the WF frames this provided a flat image, except for the bottom $\\approx 200$ pixels of each frame where the subtracted image showed some structure. \\subsection {Detection and removal of point sources} \\label{sec:lf} Once a flat image was obtained, DAOPHOT was used to find all sources greater than 4 sigma above the background and determine their magnitudes. To determine the magnitude of the point sources aperture photometry was used, with an aperture radius of $0.''5$. Crowding was not a problem on any of the frames. After zero-pointing the magnitudes of the point sources using the method described in Holtzman \\etal (1995), a point source number count was then constructed using 0.5mag bins down to the completeness limit of the data. The sum of a Gaussian globular cluster luminosity function (GCLF) \\cite{Harris88} and a power-law number count for background galaxies was then fitted to the observed point source number count. For the GCLF $\\sigma = 1.4$ mag was adopted and $M^V_0 = -7.4$ \\cite{Harris88}. For the background galaxies the slope and zero-point of the number count was found by combining I band counts from a variety of sources \\cite{Tyson88,Lilly91,Driver94,HM84,Koo86,Glaze95,Smail96,Driver95} and then fitting a power law to the combined counts (see Figure~\\ref{fi:backgals}). This gave a slope of $0.322\\pm0.005$ for the I band counts, with an intercept of $-3.158 \\pm 0.122$. The galaxy number count was fixed, as was the apparent magnitude of the GCLF, while the amplitude of the GCLF was allowed to vary to provide the best fit. The apparent magnitude of the peak of the GCLF was first set by the distances given in Ciardullo \\etal (1993) for each galaxy, then determining the distance to the galaxy using the HST data, recalculating the peak based on the new distances, and then iterating. \\begin{figure*} \\epsfig{file=backgals.eps,height=10.0cm} \\caption{Galaxy counts taken from a range of sources. All magnitudes have been converted to $I_{KC}$ and have been de-reddened. The straight line is the best least squares fit to all of the data points. The dashed lines show the error limit of the fit. } \\label{fi:backgals} \\end{figure*} Once the best fit to the point source number count was determined it was then possible to calculate the contribution to the variance in the image that point sources beyond our cutoff magnitude ($I_{KC}\\sim$24mag) make. This was done by integrating the fitted number count past the cutoff, using the equation for the residual power $P_r$ \\begin{equation} P_r = \\sum n(m)f^2(m)/\\overline{g} \\label{eqn:pr} \\end{equation} \\noindent where $n(m)$ is number of undetected point sources per pixel in the magnitude bin centred on $m$, $f(m)$ is the flux of a source of magnitude $m$ and $\\overline{g}$ is the average galaxy flux per pixel. \\subsection{Determining the SBF power by Fourier analysis} To determine the power present in the image due to SBF signal it is necessary to construct a Fourier power spectrum of the image. Since the SBF signal has the effect of the point spread function (PSF) of the telescope/detector imprinted on it, it is possible to separate it from the other sources of variance (shot noise, readout noise) that appear as white noise in the power spectrum (\\ie equal power at all scales). Before the power spectrum is found for the image, a mask image is constructed. The mask image is 1 everywhere except for areas of 20 by 20 pixels, which are set to zero, centred on the positions of all point sources down to the magnitude limit used above to calculate the number count. Also, any defects and, in the case of the PC chips, the central regions of the galaxy were masked out. The image was then multiplied by the mask, and the observed Fourier power spectrum ($P(k)$) calculated for the masked image, using the {\\it powerspec} task available under STSDAS/IRAF. The output power spectrum was normalized by the mean galaxy flux of the unmasked pixels in the image. An expectation power spectrum $E(k)$ was constructed for each region by convolving the power spectrum of the PSF with the power spectrum of the mask for that region. The PSF for each region was determined by using a grid of artificial PSFs provided by Tanvir (private communication), who used the {\\it Tiny Tim} \\cite{Kirst94} program to calculate artificial PSFs. In regions where bright stars were present the power spectra of the stars were found and compared to the power spectrum from the artificial PSFs. The agreement was found to be extremely good. The expectation power spectrum were normalised so that $E(k=0) = 1$. For each of the images under study the observed power spectrum was fitted using the relationship \\begin{equation} P(k) = P_0 E(k) + P_1 \\label{equation1} \\end{equation} \\noindent where $P_0$ is the sum of the the power in the image due to SBF and residual point sources, \\ie $P_0 = P_{fluc} + P_r$ To be able to fit the data using Equation~(\\ref{equation1}), the two dimensional power spectra (as output from the 2D Fourier analysis) were converted into one dimensional power spectra by azimuthally averaging each region --- each region was deliberately made square to make this task easier. The expectation power spectrum was then fitted to $P(k)$ using an iterative least squares method over several ranges of wave numbers (typically $40 5500$~K). In the same graph we have plotted the Li$_{\\rm max}$ values obtained for 12 Gyr, $\\alpha$ = 1.6 and [Fe/H] = -2. The initial value is taken as A(Li) = 2.35 (where A(Li) represents the logarithm of the lithium abundance in the log H =12 scale). The comparison of the Li$_{\\rm max}$ curve with the observed lithium ``plateau\" allows to determine the primordial lithium value. \\placefigure{fig3} \\placetable{table3} \\placetable{table4} Table 3 gives the values of A(Li) as obtained without any correction from the analytical fits of BM97 (their table 2). These values are precisely computed for the effective temperatures given by our models for three masses and three ages. As a further step, they are then increased by the logarithm of the corresponding Li$_{\\rm max}$/Li$_0$ value (given in Table 1). The new values are labelled as A(Li$_0$) in Table 4. If the observed lithium was exactly Li$_{\\rm max}$, all the A(Li$_0$) values should be identical. In reality, we expect the 0.70~M$_{\\odot}$ star to have suffer extra nuclear depletion, as the bottom of the convection zone is very close to the nuclear destruction layer. Also the age of these Pop II stars is supposedly larger than 10 Gyr. For these reasons, we chose to compare the four values given by the 0.75 and 0.80~M$_{\\odot}$ models at 12 and 14 Gyr. Considering an uncertainty of about 10 $\\%$ on the surface lithium abundance compared to Li$_{\\rm max}$, and taking into account the systematic error in the observations quoted by Bonifacio \\& Molaro (1997), we give for the primordial lithium abundance: A(Li$_0) = 2.35 \\pm 0.10$. When compared to BBN computations (e.g. Copi, Schramm \\& Turner 1995) this result leads to a baryonic number between 1.2 and 5 $10^{-10}$. For H~=~50, this value corresponds to $0.018 < \\Omega _{b} < 0.075$. \\vspace{2cm} We thank R.Cayrel for fruitfull comments on the manuscript. We are grateful to the Centre National Universitaire Sud de Calcul for providing computer facilities for this study. We thank the Institute for Nuclear Theory at the University of Washington where this work was initiated. C.C. acknowledges support provided by the Space Telescope Science Institute." + }, + "9802/astro-ph9802065_arXiv.txt": { + "abstract": "The low-mass X-ray binary pulsar \\stts\\ shows 0.048 Hz quasi-periodic oscillations (QPOs) and red noise variability as well as coherent pulsations at the 0.130 Hz neutron star spin frequency. Power density spectra of observations made with the {\\it Rossi X-ray Timing Explorer\\/} show significant sidebands separated from the pulsar spin frequency (and its harmonics) by the QPO frequency. These show that the instantaneous amplitude of the coherent pulsations is modulated by the amplitude of the QPOs. This phenomenon is expected in models such as the magnetospheric beat frequency model where the QPOs originate near the polar caps of the neutron star. In the 4--8 keV energy range, however, the lower-frequency sidebands are significantly stronger than their higher-frequency complements; this is inconsistent with the magnetospheric beat frequency model. We suggest that the 0.048 Hz QPOs are instead produced by a structure orbiting the neutron star at the QPO frequency. This structure crosses the line of sight once per orbit and attenuates the pulsar beam, producing the symmetric (amplitude modulation) sidebands. It also reprocesses the pulsar beam at the beat frequencies between the neutron star spin frequency and the QPOs, producing the excess variability observed in the lower-frequency sidebands. Quite independently, we find no evidence that the red noise variability modulates the amplitude of the coherent pulsations. This is also in contrast to the expectations of the magnetospheric beat frequency model and differs from the behavior in some high-mass X-ray binary pulsars. ", + "introduction": "\\label{sec:intro} Quasi-periodic oscillations (QPOs) have been detected in at least 9 accretion-powered binary X-ray pulsars (see \\cite{Take94} and references therein; also \\cite{Finger96}; \\cite{Belloni90}; \\cite{Zhang96}; \\cite{Kommers97}). Several physical mechanisms have been proposed to explain the QPOs in X-ray binaries (see \\cite{Lewin88}; \\cite{vdKlis95} for reviews). One well-developed model is the magnetospheric beat frequency model (MBFM) in which the QPO centroid frequency ($\\nu^{MBFM}_{QPO}$) represents the beat frequency between the Keplerian orbital frequency at the inner edge of the accretion disk ($\\nu_K$) and the neutron star spin frequency ($\\nu_s$): $\\nu^{MBFM}_{QPO} = \\nu_K - \\nu_s$ (\\cite{Alpar85}; \\cite{Lamb85}). Another model, which has received less discussion in the literature, is the Keplerian frequency model (KFM; \\cite{Bath74}; \\cite{vdKlis87}). In the KFM, inhomogeneities in the plasma orbiting at the inner edge of the accretion disk modulate the X-ray intensity by varying the optical depth along the line of sight; the QPO frequency equals the Keplerian frequency at the inner edge of the disk ($\\nu^{KFM}_{QPO} = \\nu_K$). \\stts\\ is a low-mass X-ray binary (LMXB) pulsar with a 7.67 s pulse period (130 mHz spin frequency). Despite extensive searches, no Doppler shifts have been detected in the X-ray pulse arrival times (\\cite{Levine88}; \\cite{Shinoda90}; \\cite{Chakra98}). The presence of a low-mass companion star in a 42 minute prograde orbit around the neutron star is inferred from photometric timing measurements on the optical counterpart, KZ TrA, which shows pulsations at the same frequency as the X-ray pulsations. The power spectrum of the optical variability shows additional weak pulsations in a sideband 0.4 mHz below the main optical pulsation frequency (\\cite{Middle81}; \\cite{Chakra97}). The interpretation is that the sideband represents X-rays from the pulsar beam that have been reprocessed on the surface of the companion star. The 0.4 mHz shift to lower frequencies results from the lower apparent pulsar frequency observed in a frame rotating with the binary orbit of the companion (\\cite{Middle81}). The long-term behavior of the X-ray pulse frequency has been discussed by Chakrabarty et al.\\ (1997). Although the pulsar was spinning up ($\\dot{\\nu_s} > 0$) for more than 13 years after its discovery, it experienced a torque reversal in approximately June 1990. It has been spinning down ($\\dot{\\nu_s} < 0$) ever since (\\cite{Chakra96}). QPOs at 48 mHz with fractional root-mean-square (rms) amplitudes as high as 17\\% have been seen in X-ray observations with {\\it Ginga\\/} (\\cite{Shinoda90}), {\\it ASCA\\/} (\\cite{Angelini95}), and {\\it RXTE\\/} (\\cite{Chakra97}). Chakrabarty (1998) recently detected these QPOs in the optical $U$, $B$, and $R$ bands. In this {\\it Letter\\/} we present evidence that neither the MBFM nor the KFM provides a suitable explanation for the 48 mHz QPOs in \\stts. We propose instead that the QPOs are produced by some kind of coherent structure (a ``blob'') orbiting the neutron star with an orbital frequency equal to the QPO frequency, but the ``blob'' is not at the inner edge of the accretion disk. ", + "conclusions": "\\label{sec:concl} Significant sidebands are detected around the 130 mHz pulsar frequency and its harmonics in power density spectra of 4U~1626$-$67. The sidebands appear at frequencies $n \\nu_s \\pm \\nu_{QPO}$, where $n = 1,2, \\ldots$ is an integer. In the 17--30 keV range the sidebands are mirror images of each other: they are symmetric in both frequency and power amplitude. We will refer to these as ``symmetric'' side bands. In the 4--8 keV range, however, the lower-frequency sidebands contain significantly more power than the higher-frequency ones: they are symmetric in frequency but {\\it not} in power amplitude. We will assume that these ``asymmetric'' sidebands represent the superposition of underlying symmetric sidebands {\\it plus\\/} some additional power at the lower sideband frequencies ($n \\nu_s - \\nu_{QPO}$). We will refer to this excess power as ``enhanced lower-frequency'' sidebands. The presence of the symmetric sidebands suggests that the instantaneous amplitude of the coherent pulsations contains a term proportional to the 48 mHz QPO signal. This phenomenon is expected in models where some of the QPOs are produced near the polar caps of the neutron star. In the magnetospheric beat frequency model (MBFM), for example, magnetically gated clumps of matter from the inner accretion disk would follow the magnetic field lines onto the polar caps and modulate the X-ray intensity there (\\cite{Alpar85}; \\cite{Lamb85}). Symmetric sidebands could also occur in models such as the Keplerian frequency model (KFM) where inhomogeneities in the accretion disk quasi-periodically absorb some of the pulsar beam along the line of sight. The presence of the enhanced lower-frequency sidebands shows that the situation is more complicated, however. In the 4--8 keV range at least two of the lower-frequency sidebands (the one at $\\nu_s - \\nu_{QPO} = 82$ mHz and the one at $2 \\nu_s - \\nu_{QPO} = 212$ mHz) are much stronger than their higher-frequency complements. Since amplitude modulation produces symmetric sidebands, there must be a physical mechanism that produces additional variability at the frequencies $n \\nu_s - \\nu_{QPO}$. Neither the MBFM nor the KFM (as usually formulated) provides a satisfactory explanation for both the main QPO signal and the enhanced lower-frequency sidebands. In the MBFM the QPOs represent a beat frequency between the Keplerian frequency at the magnetopause and the pulsar frequency ($\\nu_{QPO}^{MBFM} = \\nu_{K} - \\nu_{s}$). If we set $\\nu_K = 178$ mHz, the MBFM would predict the 48 mHz QPOs. The presence of symmetric sidebands would follow since the 48 mHz oscillations occur in the accretion stream near the neutron star's polar caps. But this MBFM provides no mechanism that allows the QPOs to again ``beat'' against the rotating pulsar beam to produce the enhanced lower-frequency sidebands at $\\nu_s - \\nu_{QPO} = 82$ mHz and $2 \\nu_s - \\nu_{QPO} = 212$ mHz (\\cite{Alpar85}; \\cite{Lamb85}; \\cite{Shib87}). The KFM can explain the sideband structure (see our modified version of this model below) but not the direct QPOs. This is because the KFM cannot explain QPOs with centroid frequencies {\\it less\\/} than the neutron star spin frequency. In the KFM the QPO frequency would be the Keplerian frequency at the inner edge of the accretion disk (Bath et al.\\ 1974; \\cite{vdKlis87}). Since the QPO centroid frequency is less than the pulsar spin frequency in \\stts, the inner edge of the disk would lie outside the co-rotation radius. Centrifugal forces exerted by the rotating neutron star magnetosphere would therefore inhibit accretion (the ``propeller effect''; \\cite{Ill75}). The radius corresponding to a 48 mHz Keplerian orbital frequency is $r_K = 1.3 \\times 10^9$ cm, which exceeds the co-rotation radius $r_{co} = 6.5 \\times 10^{8}$ cm. The location of the inner edge of the disk is less certain: depending on the distance to the source, $r_K$ could exceed the Alfv\\'en radius, $r_A \\approx 3 \\times 10^{9} d_{{\\rm kpc}}^{-4/7}$ cm $\\approx 9 \\times 10^{8}$ cm, where we have assumed a distance of 8 kpc and used the measurements by Orlandini et al. (1997) of the 0.1--100 keV luminosity ($L_X = 6.6 \\times 10^{34}$ erg s$^{-1}$ $d_{{\\rm kpc}}^2$) and magnetic field strength ($B = 3.3 \\times 10^{12}$ G). If we dispense with the KFM's requirement that the QPO frequency represents the {\\it Keplerian frequency at the inner edge of the disk}, however, we obtain a scenario that can simultaneously account for the direct QPOs, the symmetric sidebands, and the enhanced lower-frequency sidebands. Suppose a large coherent structure (a ``blob'') of material orbits the neutron star with a frequency roughly equal to the QPO centroid frequency (48 mHz); this may or may not represent the Keplerian frequency at the radius of the blob. We use the term ``blob'' to distinguish this structure from the ``clumps'' of the MBFM (\\cite{Lamb85}) and the ``inhomogeneities'' of the KFM (Bath et al.\\ 1974). The direct QPO signal is produced as the ``blob'' modulates the optical depth to the accretion disk as it orbits. Once every orbit a portion of the ``blob'' crosses the line of sight between the neutron star and the Earth and scatters X-rays from the pulsar beam out of the line of sight. This quasi-periodic attenuation of the pulsar beam intensity produces the symmetric sidebands around the spin frequency and its harmonics. Suppose also that the ``blob'' orbits in the same sense as the pulsar rotation, so it is illuminated by the pulsar beam with a frequency of $\\nu_s - \\nu_{QPO} = 82$ mHz. The first harmonic illuminates the ``blob'' with a frequency of $2\\nu_s - \\nu_{QPO} = 212$ mHz. When the ``blob'' is not crossing the line of sight between the neutron star and the Earth, it reprocesses X-rays from the pulsar beam and returns some of them along the line of sight. The reprocessed radiation appears as oscillations at the frequencies of the enhanced lower-frequency sidebands, $n \\nu_s - \\nu_{QPO}$. The situation is analogous to the reprocessing of the pulsar beam by the companion star (\\cite{Middle81}; \\cite{Chakra97}). The overall scenario is similar to the production of orbital sidebands due to the modulation of pulsations in intermediate polars (\\cite{Warner86}). The fact that the enhanced lower-frequency sidebands are found only below the fundamental (130 mHz) and the first harmonic (260 mHz) in the 4--8 keV range needs an explanation. We suggest that this is related to the reprocessing of the pulsar beam. In the 17--30 keV range the pulse profile is dominated by the fundamental and the first harmonic (see Figure \\ref{fig:fullps}). If the ``blob'' reprocesses hard X-rays and emits them at lower energies, the strengths of the sidebands at 82 mHz and 212 mHz in the 4--8 keV range would reflect the strengths of the corresponding pulse profile harmonics at {\\it higher\\/} energies (e.g. 17--30 keV). The nature of the reprocessing ``blob'' is not clear. It is unlikely that there are many ``blobs'' scattered around the 48 mHz orbit, because reprocessed emission from these would contribute with many different phases and tend to wash out the oscillations responsible for the enhanced lower-frequency sidebands. On the other hand, it is not clear why a single reprocessing structure of limited spatial extent would survive in the accretion disk against the differential rotation of nearby Keplerian orbits. The QPOs appear to be stable on the time scale of a decade or more, having been first observed in 1988 (\\cite{Shinoda90}). It is perhaps more likely that the orbital frequency of the ``blob'' is {\\it not} Keplerian. For example, the ``blob'' may represent a superposition of oscillatory modes traveling as a stable wave packet around the accretion disk. Alpar \\& Yilmaz (1997) have described models of the normal-branch and horizontal-branch QPOs in LMXBs in terms of wave packets of sound waves in the accretion disk. To our knowledge the sideband structure has not been previously used as a diagnostic of the QPO mechanism in X-ray pulsars. Regardless of whether the simple model we set forth above can stand up to closer scrutiny, the sideband structure provides remarkably constraining information on the possible QPO mechanisms. An unusual sideband structure was recently detected in the high-mass X-ray binary (HMXB) pulsar Cen~X-3 (M. Finger, private communication, 1997) that may require a different interpretation than the one discussed here. On the other hand, the sideband structure has been used to consider the origin of RN variability in X-ray pulsars (\\cite{Burd97}; \\cite{Lazz97}). If the RN is produced by inhomogeneities in the accretion flow onto the magnetic poles of the neutron star, then the instantaneous amplitude of the coherent pulsations should contain a term proportional to the RN signal (just as for the QPO signal discussed above). For example, in the MBFM the same magnetically gated clumps that produce the QPOs are expected to create a RN component that represents the frequency content and lifetime broadening of the shots produced by each clump (\\cite{Lamb85}; \\cite{Shib87}). The MBFM does not rule out the possibility of additional sources of RN variability that are unrelated to the magnetic gating mechanism, however. We found no detectable modulation of the pulsar beam by the RN component in \\stts\\ (see Figure \\ref{fig:fullps}) even though the QPO signal modulates the pulsar beam substantially. The fractional rms variability between 0.0025 Hz and $\\sim 1$ Hz in the combined 4--8 keV RN components was $42.5 \\pm 0.5$\\%; and that in the combined 17-30 keV RN components was $63.2 \\pm 0.6$\\%. Yet the upper limits on the total rms contained in the sidebands that would result from a modulation of the coherent pulses by the RN are only $\\sim 3$\\%. This suggests that most of RN does {\\it not\\/} originate near the polar caps. This is in contrast to case of the HMXB pulsars SMC X-1, Vela X-1, 4U~1145$-$62, and possibly Cen X-3, in which the RN variability does appear to modulate the coherent pulse amplitude (\\cite{Burd97}; \\cite{Lazz97}). To summarize, we have detected sidebands around the neutron star spin frequency and its harmonics in \\stts. These show that the QPOs modulate the amplitude of the coherent pulsations. The presence of the enhanced lower-frequency sidebands below the pulsar frequency and its harmonics is inconsistent with the MBFM. Interpreting the lower-frequency sidebands as reprocessed radiation from the pulsar beam, we have proposed a modification of previous (Keplerian frequency) obscuration models that explains the observed QPOs and sideband structure. The strong RN component below 1 Hz does not appear to significantly modulate the amplitude of the coherent pulsations." + }, + "9802/astro-ph9802253_arXiv.txt": { + "abstract": "We present high dispersion HST GHRS UV spectroscopic observations of 8 H\\,{\\sc II} galaxies covering a wide range of metallicities and physical properties. We have found Ly$\\alpha$\\ emission in 4 galaxies with blueshifted absorption features, leading to P~Cygni like profiles in 3 of them. In all these objects the O\\,{\\sc I} and Si\\,{\\sc II} absorption lines are also blueshifted with respect to the ionized gas, indicating that the neutral gas is outflowing in these galaxies with velocities up to 200~km\\thinspace s$^{-1}$. The rest of the sample shows broad damped Ly$\\alpha$\\ absorption profiles centered at the wavelength corresponding to the redshift of the H\\,{\\sc II} emitting gas. We therefore find that the velocity structure of the neutral gas in these galaxies is the driving factor that determines the detectability of Ly$\\alpha$\\ in emission. Relatively small column densities of neutral gas with even very small dust content would destroy the Ly$\\alpha$\\ emission if this gas is static with respect to the ionized region where Ly$\\alpha$\\ photons originate. The situation changes dramatically when most of the neutral gas is velocity--shifted with respect to the ionized regions because resonant scattering by neutral hydrogen will be most efficient at wavelengths shorter than the Ly$\\alpha$\\ emission, allowing the Ly$\\alpha$\\ photons to escape (at least partially). This mechanism complements the effect of porosity in the neutral interstellar medium discussed by other authors, which allows to explain the escape of Ly$\\alpha$\\ photons in regions surrounded by static neutral gas, but with only partial covering factors. The anisotropy of these gas flows and their dependence on the intrinsic properties of the violent star-forming episodes taking place in these objects (age, strength, gas geometry,...) might explain (in part) the apparent lack of correlation between other properties (like metallicity) and the frequency of occurence and strength of Ly$\\alpha$\\ emission in star-forming galaxies. Attempts to derive the comoving star--formation rate at high redshifts from Ly$\\alpha$\\ emission searches are highly questionable. ", + "introduction": "The detection of galaxies at large redshifts that are forming stars for the first time, the so--called primeval galaxies, remains a very important astrophysical challenge. Bearing in mind that galaxy formation may not be assigned to any preferential cosmological epoch but instead is probably a continuous process, one might find left--over pristine gas pockets that are forming young galaxies at the present epoch. For this reason there may be star--forming galaxies in our local universe that look very much like distant primeval ones. Hopes have been that Ly$\\alpha$\\ emission could be a signature of star formation that would be recognized up to very large redshifts; hence there have been numerous studies of the Ly$\\alpha$\\ emission from distant and local starbursts. Early IUE observations were performed on more than a dozen nearby starburst galaxies in its SWP low resolution mode (Meier \\& Terlevich 1981; Hartmann et al. 1984; Deharveng et al. 1986; Hartmann et al. 1988 and Terlevich et al. 1993). Galaxies with redshifts large enough that their Ly$\\alpha$\\ emission is separated from the geocoronal line were selected. It was realized from the very beginning that the Ly$\\alpha$/H$\\beta$ ratio and the Ly$\\alpha$\\ equivalent width are much smaller, by at least an order of magnitude, than expected from the recombination theory. These early works have also shown a possible anticorrelation between the Ly$\\alpha$/H$\\beta$ ratio and the H\\,{\\sc II} galaxy metallicity (actually the O/H abundance, as measured in the ionized gas). \\begin{figure}[t] % \\begin{center}\\mbox{\\epsfxsize=9cm \\epsfbox{h0749f2.eps}}\\end{center} \\caption[]{ Detail on the O\\,{\\sc I} and Si\\,{\\sc II} region for the galaxies showing Ly$\\alpha$\\ emission. The vertical bars indicate the wavelength at which the O\\,{\\sc I} and Si\\,{\\sc II} absorption lines should be located, according to the redshift derived from optical emission lines. Some Galactic absorption lines have been marked. Note that the metallic lines appear systematically blueshifted in these galaxies with respect to the systemic velocity. In some cases there is no significant absorption at all at zero velocity. } \\label{fig:oila} \\end{figure} \\begin{figure}[t] % \\begin{center}\\mbox{\\epsfxsize=9cm \\epsfbox{h0749f3.eps}}\\end{center} \\caption[]{ O\\,{\\sc I} and Si\\,{\\sc II} region for the galaxies showing damped Ly$\\alpha$\\ absorptions. Details as in Fig.~\\ref{fig:oila}. Note that in these galaxies the metallic lines are essentially at the same redshift than the ionized gas, indicating the presence of static clouds of neutral gas, as discussed in the text. } \\label{fig:oiab} \\end{figure} These results, the lack of ``primeval galaxies'' at large redshift in blank sky searches for redshifted Ly$\\alpha$\\ emission and the few tentative detections of Ly$\\alpha$\\ emission from the damped Ly$\\alpha$\\ systems have been attributed to the effects of dust absorption that preferentially destroys Ly$\\alpha$\\ photons (Charlot \\& Fall 1993, and references therein). The process behind this is that the transfer of Ly$\\alpha$\\ radiation is strongly affected by resonant scattering from neutral interstellar hydrogen atoms. By increasing enormously their optical path length, Ly$\\alpha$\\ photons become more vulnerable to dust absorption, even in small amounts (Neufeld 1991; Charlot \\& Fall 1991; Chen \\& Neufeld 1994). This process was believed (even in the early paper of Meier \\& Terlevich 1981) to be able to account for the anticorrelation between the Ly$\\alpha$\\ emission line visibility and the dust abundance in these galaxies, as parameterized by the metallicity. Alternatively such an anticorrelation has been attributed to a metallicity--dependent extinction law at the wavelength of Ly$\\alpha$\\ (Calzetti \\& Kinney 1992, Valls-Gabaud 1993). However this conclusion seems very unlikely in view of the anticorrelation between the Ly$\\alpha$\\ equivalent width and the gas-phase abundance of oxygen O/H. Charlot \\& Fall (1993) have emphasized the advantages of using the Ly$\\alpha$\\ equivalent widths rather than the Ly$\\alpha$/H$\\beta$ ratios, because the former are independent on the extinction curve of the dust and can be measured with a single observational device. Their discussion of the anticorrelation between the Ly$\\alpha$\\ line equivalent widths and the O/H abundances in a sample of nearby star-forming galaxies has examined several factors that will affect the observed Ly$\\alpha$\\ emission from galaxies, among which contributions from supernova remnants and active galactic nuclei, the orientation of the galaxy and the absorption by dust. They finally suggest that the structure of the interstellar medium (porosity and multi-phase structure of the medium) is most probably the dominant one. \\begin{table*}[t] \\caption{Adopted properties of observed H\\,{\\sc II} galaxies. Systemic velocities have been taken from the NASA Extragalactic Database, except for Haro~2 (Legrand et al. 1997). } \\label{tab:galaxies} \\begin{tabular}{lcccc}\\hline Galaxies & m(V or B) & v(km\\thinspace s$^{-1}$) & 12+log(O/H) & E(B-V) \\\\ \\hline ESO 350-IG038 & 14.27V & 6156 & ?? & 0.16\\\\ SBS0335-052 & 16.65V & 4043 & 7.36 & 0.18 \\\\ IRAS 08339+6517& 14.16V & 5730 & ?? & 0.55 \\\\ IZw 18 & 15.6B & 740 & 7.17 & $<$0.10 \\\\ Haro 2 & 13.4V & 1465 & 8.40 & 0.12 \\\\ Mkn 36 & 15.5V & 646 & 7.86 & $<$0.10 \\\\ IIZw 70 & 14.83V & 1215 & 8.33 & 0.15 \\\\ ESO 400-G043 & 14.22B & 5900 & 8.0 & 0.20 \\\\ \\hline \\end{tabular} \\end{table*} Our new HST observations indicate that velocity structure in the interstellar medium plays a key role in the transfer and escape of Ly$\\alpha$\\ photons. At first place, Ly$\\alpha$\\ was observed only in absorption in the starburst dwarf galaxy IZw~18 by Kunth et al. (1994). Since IZw~18 at $Z=1/50 ~Z_{\\odot}$ is the most metal--poor starburst galaxy known at present, it was considered previously a good candidate to show Ly$\\alpha$\\ in emission. To add to the confusion, a positive Ly$\\alpha$\\ emission showing a complicated profile, but a clear P~Cygni component, has been detected in Haro~2, a rather dusty star-forming galaxy at $Z=1/3 ~Z_{\\odot}$ (Lequeux et al. 1995). Giavalisco et al. (1996) have strengthened the suggestion that the transport of the Ly$\\alpha$\\ photons is primarily controlled by the ISM geometry rather than by the amount of dust, so that the Ly$\\alpha$\\ emission line would be detected only if there are holes (regions with low column density of neutral gas) along the line of sight, a factor which in principle is independent on dust and metal content of the gas. As we show hereafter, other factors can be certainly more important in accounting for variations in the Ly$\\alpha$\\ emission strength, at least in some cases. The detection of a P~Cygni profile in the Ly$\\alpha$\\ emission line of Haro~2 led us to postulate that the line was visible because the absorbing neutral gas was velocity--shifted with respect to the ionized gas. This was confirmed by the analysis of the UV O\\,{\\sc I} and Si\\,{\\sc II} absorption lines, which were blue--shifted by 200~km\\thinspace s$^{-1}$\\ with respect to the optical emission lines, and also of the profile of the H$\\alpha$ line (Legrand et al. 1997). These new facts and the capability of the HST to analyze in detail for the first time Ly$\\alpha$\\ line profiles in nearby galaxies led us to embark on a longer--term project using the GHRS aiming to study the processes controlling the detectability of the Ly$\\alpha$\\ emission line in star-forming galaxies. These studies have also been aimed to measure abundances in the neutral gas of gas--rich dwarf galaxies with spectra dominated by recent star formation episodes. Indeed, in objects such as these, the H\\,{\\sc I} clouds largely extend beyond the optical images suggesting that a substantial fraction of this gas might still be chemically unevolved or even pristine (Roy \\& Kunth 1995). At a spectral resolution of 20 000 it became possible to disentangle nebular from stellar absorption lines and to give crude estimates of the metal abundances in the interstellar medium. The study of the IZw~18 data by Kunth et al. (1994) and the preliminary analysis of the rest of the sample (Kunth et al. 1997) have yielded extremely low values of the O\\,{\\sc I}/H\\,{\\sc I} ratios (log N(O\\, {\\sc I})/N(H\\,{\\sc I}) $<$ -7) in some galaxies of the sample. The complete analysis of the interstellar abundances will be presented in a forthcoming paper. In Sect.~2 we present new HST data on a sample of 8 H\\,{\\sc II} galaxies. Spectra are described in Sect.~3 and the results are discussed in Sect.~4. The conclusions are finally summarized in Sect.~5. ", + "conclusions": "We have analyzed HST UV spectroscopical data of eight H\\,{\\sc II} galaxies aiming to characterize the detectability of the Ly$\\alpha$\\ emission line in this kind of objects. We obtain the following results: \\begin{itemize} \\item Ly$\\alpha$\\ emission has been observed in four out of the eight H\\,{\\sc II} galaxies. In all these four galaxies we have found a clear evidence of a wide velocity field by the presence of deep absorption troughs at the blue side of the Ly$\\alpha$\\ profiles. Moreover, absorption lines of metallic elements (O\\,{\\sc I}, Si\\,{\\sc II}) are also significantly blueshifted with respect to the H\\,{\\sc II} gas velocity. \\item The determining factor for the detectability of the Ly$\\alpha$\\ emission line in these galaxies is therefore the velocity structure of the neutral gas along the line of sight, rather than the abundance of dust particles alone. If most of the neutral gas is outflowing from the ionized region, the Ly$\\alpha$\\ emission line would escape (partially) unaffected, independently on the metal abundance and dust content of this neutral gas. This outflowing material apparently powered by massive stars winds and/or SN may eventually leave the galaxy. We thus may be witnessing galactic winds resulting from intense star formation activity. In the case of Haro~2, Lequeux et al. (1995) suggested that 10$^{7}$ M$\\odot$ are expanding at 200 km\\thinspace s$^{-1}$. \\item Broad Ly$\\alpha$\\ absorption is detected in all H\\,{\\sc II} galaxies. The derived N(H\\,{\\sc I}) column densities lie unexpectedly inside a relatively small range with 6 of the 8 H\\,{\\sc II} galaxies having logarithmic column densities logN(H\\,{\\sc I}) between 19.9 and 21.1~cm$^{-2}$\\ (extreme values are 19.7 and 21.5). We stress again that the Ly$\\alpha$\\ photons emitted by the H\\,{\\sc II} region are absorbed or redistributed by the H\\,{\\sc I} gas only if its velocity is the same as that of the H\\,{\\sc II} region. Otherwise, the photons that are resonantly trapped were emitted in the stellar continuum close to Ly$\\alpha$. \\item The dependence of Ly$\\alpha$\\ emission detectability on the presence/absence of neutral static/outflowing gas along the line of sight (and within the field of view covered by the slit), helps to explain the apparently contradictory detection of Ly$\\alpha$\\ emission in metal and dust--rich galaxies (like Haro~2), while it may be absent in metal and dust deficient objects, of which IZw~18 is the prototype. \\item A partial covering factor of the H\\,{\\sc II} region by neutral gas, with low H\\,{\\sc I} column densities, would be required to allow the detection of the Ly$\\alpha$\\ emission line if the neutral gas is static with respect to the ionized regions. \\item The generally weak or absent Ly$\\alpha$\\ emission from ``primeval'' and other galaxies at high redshifts can only be explained by velocity-structure effects combined with absorption of the Ly$\\alpha$\\ photons by dust grains. The relatively small angular extent of these sources implies that if photons were leaking through the neutral gas clouds surface after multiple scattering without being destroyed, the equivalent widths of the lines measured from Earth should be significantly higher than observed. \\item The present study invalidates attempts to measure the comoving star--formation rate density at high redshift on the basis of Ly$\\alpha$\\ emission surveys. \\end{itemize} Future work should address the several effects discussed in this work to understand the reasons that govern the presence/absence of the Ly$\\alpha$\\ line emission and absorption: the strength and age of the burst, the metallicity of the gas (controlling the cooling, hence the wind evolution), the gravitational potential of the parent galaxy and its morphology and the H\\,{\\sc I} and the dust distributions will all play a role. The challenge is to determine their relative importance in affecting the Ly$\\alpha$\\ emission and absorption processes. Clearly the way forward is to realistically model the hydrodynamical evolution of the ISM in gas rich dwarf galaxies under the influence of starburst of different fractional masses. Particular attention should be paid to the time evolution of neutral gas kinematical and structural parameters." + }, + "9802/hep-ph9802351_arXiv.txt": { + "abstract": "The current status of Ultra High Energy Cosmic Rays (UHECR) is reviewed, with emphasis given to theoretical interpretation of the observed events. The galactic and extragalactic origin, in case of astrophysical sources of UHE particles, have the problems either with acceleration to the observed energies or with the fluxes and spectra. Topological defects can naturally produce particles with energies as observed and much higher, but in most cases fail to produce the observed fluxes. Cosmic necklaces and monopole-antimonopole pairs are identified as most plausible sources, which can provide the observed flux and spectrum. The relic superheavy particles are shown to be clustering in the Galactic halo, producing UHECR without Greisen-Zatsepin-Kuzmin cutoff. The Lightest Supersymmetric Particles are discussed as UHE carriers in the Universe. ", + "introduction": "Cosmic rays (CR) are observed in a wide energy range, starting from subGeV energies and up to $3\\cdot 10^{20}~eV$ (see Fig.1). Apart from the highest energies, these particles are accelerated in our Galaxy, most probably, by shocks produced by SN II explosions. Up to energy $10^{15} - 10^{16}~eV$ the CR flux is dominated by protons, at higher energies CR have the mixed composition, and there are indications that at energies about $\\sim 10^{17}~eV$ iron nuclei dominate in the CR flux. In a wide range of energies from $1~GeV$ up the $3\\cdot 10^{15}~eV$ the spectrum is power-law $ \\sim E^{-2.65}$, at energy $ 3\\cdot 10^{15}~eV$ the spectrum steepens and becomes $\\sim E^{-3.1}$ at $E > 10^{17}~eV$. At $E \\geq 10^{19}~eV$ a new more flat component appears (see Fig.1). The highest energies detected so far are $2-3\\cdot 10^{20}~eV$. The first steepening of the spectrum (the knee) at energy $3\\cdot 10^{15}~eV - 1\\cdot 10^{16}~eV$ is usually explained by inefficient confinement of CR in the Galaxy. This process must be accompanied by enrichment of heavy nuclei in CR composition at energy $\\sim 10^{15}~eV$ and above. There are some indications that such enrichment is really observed. There is no universal definition for Ultra High Energy Cosmic Rays (UHECR). Sometimes this term is applied for $E > 1\\cdot 10^{17}~eV$ or $E> 1\\cdot 10^{18}~eV$. I shall use this term for $E > 1\\cdot 10^{19}~eV$, where the new flat component appears. It is natural to think that this component has extragalactic origin, though, in principle, very large halo with regular magnetic field can confine particles of these energies, especially if they are heavy nuclei. UHECR of extragalactic origin have a signature called the Greisen-Zatsepin-Kuzmin (GZK) cutoff \\cite{GZK}. This phenomenon is caused by energy losses of UHE protons due to pion production in collisions with microwave photons. The energy losses start sharply increasing at $E \\sim 3\\cdot 10^{19}~eV$ (Fig.2). This energy is connected with energy of the spectrum steepening (\"cutoff\") in the model-dependent way. In case the sources are distributed uniformly in the Universe (standard assumption), the steepening starts at $E_{bb} \\approx 3\\cdot 10^{19}~eV$. The flux at $E>E_{bb}$ is produced by nearby sources. If there is a local enhancement of the sources, $E_{bb}$ increases \\cite{BBDGP};in case the sources are located at large distances, $E_{bb}$ decreases and steepening is exponential. It is more convenient to characterize steepening by energy $E_{1/2}$ \\cite{BBDGP}, where the flux becomes half of the power-law extrapolation of unmodified flux. In case of uniform distribution of the sources $E_{1/2} \\approx 5.8\\cdot 10^{19}~eV$ \\cite{BBDGP} for a wide range of exponents $\\gamma$ of generation spectrum. Apart from GZK cutoff, there may be two more signatures of extragalactic cosmic rays: a bump and a dip in differential spectrum which precede the cutoff \\cite{HS85,BG87}. The bump is a consequence of a number conservation of protons in the spectrum: protons loose energy and are accumulated before the cutoff. The dip is formed due to pair-production ($e^+e^-$) energy losses of UHE proton. The both features show up most clearly in the differential spectrum of a single distant source in the case of a flat generation spectrum. In diffuse spectra (from many sources) these features are weak or absent. \\noindent {\\em UHE nuclei} spectra exhibit steepening (\"cutoff\") approximately at the same energy as protons, though due to different physical processes (see \\cite{BBDGP} for a review). The relevant energy losses are caused by photodisintegration of nuclei at collisions with microwave photons, and the steepening energy is determined by energy, when photodisintegration energy-losses start to dominate over adiabatic ones (Fig.2.). {\\em UHE photons} with $E_{\\gamma} \\sim 10^{19} - 10^{22}~eV$ have an absorption length less than $10~Mpc$, mainly due to interaction with radio-background \\cite{B70,PB}. The observation of cosmic ray particles with energies higher than $10^{20}~eV$ gives a serious challenge to our understanding of origin of UHECR: What are the mechanisms of acceleration? Why the GZK cutoff is absent? ", + "conclusions": "At $E\\geq 1\\cdot 10^{19}~eV$ a new component of cosmic rays with a flat spectrum is observed. Two highest energy events have $E \\approx 2-3\\cdot 10^{20}~eV$. According to the Fly's Eye and Yakutsk data the chemical composition is better described by protons than heavy nuclei. The AGASA data are consistent with isotropy in arrival of the particles, though theoretical analysis reveals some correlation of arrival direction with Local Supercluster plane. AGASA has observed clustering of UHE events: three pairs of particles with small angular separation. The galactic origin of UHECR is disfavored: the maximal observed energies are less than that known for the galactic sources, and the strong Galactic disc anisotropy is predicted even for the extreme magnetic fields in the disc and halo. The signature of extragalactic UHECR is GZK cutoff. The position of steepening is model-dependent value. For the Universe uniformly filled with sources, the steepening starts at $E_{bb} \\approx 3\\cdot 10^{19}~eV$ and has $E_{1/2} \\approx 6\\cdot 10^{19}~eV$ (the energy at which spectrum becomes a factor of two lower than a power-law extrapolation from lower energies). The spectra of UHE nuclei exhibit steepening approximately at the same energy as protons. UHE photons have small absorption length due to interaction with radio background radiation. The extragalactic astrophysical sources theoretically studied so far, have either too small $E_{max}$ or are located too far away. The Local Supercluster (LS) model can give spectrum with $E_{1/2} \\sim 10^{20}~eV$, if density contrast for the sources (the ratio of densities inside LS and outside) is larger than 10. Topological Defects naturally produce particles with extremely high energies, much in excess of what is presently observed. However, the fluxes from most known TD are too small. So far only necklaces and monopole-antimonopole pairs can provide the observed flux of UHE CR. Another promising sources of UHE CR are relic superheavy particles. These particles should be clustering in the halo of our Galaxy, and thus UHECR produced at their decays do not have the GZK cutoff. The signatures of this model are dominance of photons in the primary flux and Virgo cluster as a possible discrete source. Apart from protons and photons, the light gluinos can be successful UHE carriers, but they are disfavored in mass interval at interest." + }, + "9802/astro-ph9802208_arXiv.txt": { + "abstract": "\\vskip .2 truein Analytic expressions for distance--redshift relations which have been corrected for the effects of inhomogeneities in the Friedmann-Lema\\^itre-Robertson-Walker (FLRW) mass density are given in terms of Heun functions and are used to illustrate the significance of inhomogeneities on a determination of the mass parameter $\\OM$ and the cosmological constant $\\Lambda$. The values of these parameters inferred from a given set of observations depend on the fractional amount of matter in inhomogeneities and can significantly differ from those obtained by using the standard magnitude-redshift ($m$-$z$) result for pure dust FLRW models. As an example a determination of $\\OM$ made by applying the homogeneous distance--redshift relation to SN 1997ap at $z=0.83$ could be as much as 50\\% lower than its true value. ", + "introduction": "\\label{sec-intro} When attempting to evaluate the mass parameter $\\OM$ and/or the cosmological constant $\\Lambda$, observations of quantities such as magnitude, angular separation, and redshift are made on objects distant enough for curvature effects to be detected. As an example, for Type Ia supernovae (SNe Ia) corrected magnitudes and redshifts ($m$-$z$) are measured, plotted, and compared with theoretical $m(\\OM,\\Lambda; z)$ curves computed for the FLRW models (\\cite{PS1}, \\cite{PS2}, \\cite{GP}). In spite of the fact that the FLRW models contain only homogeneously and isotropically distributed perfect fluid gravity sources, one of these models is assumed to represent the ``large scale\" geometry of the universe. Relations like $m(\\OM,\\Lambda;z)$ are also commonly assumed to be valid, on average. This latter assumption may well be incorrect for some distant observations including SNe Ia, but even if technically correct may not be useful in determining $\\OM$ and $\\Lambda$. In particular if the underlying mass density approximately follows luminous matter (\\ie associated with bounded galaxies) then effects of inhomogeneities on relations like $m(\\OM,\\Lambda;z)$ must be taken into account. The majority of currently observed SNe Ia are not being seen through foreground galaxies and whether or not this is due entirely to selection (rather than statistics) is not important. If the objects observed do not have the average FLRW mass density $\\rho_0$ in their foregrounds then the FLRW \\mz\\ relation does not apply to them (see \\cite{KR2}). Ultimately some SNe Ia should exist behind foreground galaxies (\\cite{RK}) and for these, \\mz\\ should be computed using the lensing formulas. These formulas (\\cite{BR} and \\cite{CJ}) contain source-observer, deflector-observer, and source-deflector distances, respectively $D_s, D_d$, and $D_{ds}$, all of which depend on the mass density in the observing beam, \\underline{excluding} the deflector. These distances will not be given by the standard FLRW result if the observing beam contains less than the average FLRW mass density. In \\S \\ref{sec-optics} the average area-redshift equation (\\ref{Af}) for a light beam traveling through a FLRW Swiss cheese universe is given and its solution is related to the luminosity distance--redshift relation $D_{\\ell}(z)$. In \\S \\ref{sec-B=0} the solution of this equation is given for the case where gravitational lensing can be neglected. The new result of this paper, $D_{\\ell}(z)$ without lensing for FLRW Swiss cheese can be found in equations (\\ref{Dell0}) and (\\ref{Dellinfty}), and for the special case $\\OO=1$ in equations (\\ref{Dell01}) and (\\ref{Dellinfty1}) of Appendix A. In \\S \\ref{sec-mzplots} numerous \\mz plots are given to illustrate the effects of inhomogeneities and some conclusions are drawn. It is argued that if homogeneities are not taken into account when attempting to determine $\\Omega_m$ and $\\Lambda$, errors as large as 50\\% could be made. Even though the $D_{\\ell}(z)$ given here has been derived using the exact Swiss cheese cosmologies, the result are valid for observations in essentially any perturbed pressure-free FLRW models in which lensing can be neglected. Inhomogeneous models of the Swiss cheese type and their associated optical equations discussed here are often mistakenly attributed to Dyer and Roeder (see Appendix B). Appendix C contains some useful simplifications for evaluating the real-valued Heun functions needed in the analytic \\mz relations given here. Appendix C also contains six useful lines of Mathematica code which numerically evaluates and plot these same \\mz relations. ", + "conclusions": "\\label{sec-mzplots} In this section several magnitude-redshift plots are given to illustrate the effects that density clumps can have on the \\mz relation and consequently on a determination of $\\OM$ and $\\Lambda$ made by using this relation. Because \\mz depends differently on $\\OM$ and $\\OL$ as a function of redshift for the FLRW models, both parameters could in principle be determined from a sufficient quantity of accurate SNe Ia data. Clumping provides an additional parameter $\\nu$ which complicates any such determination. As can be seen from (\\ref{Dlseries}) the dependence of \\mz on this additional parameter could also be determined by enough data. However, such a triple determination is certainly more complicated. What will be done here to illustrate the effects of the $\\nu$ parameter is to plot multiple \\mz curves for various values of all three parameters $\\nu, \\ \\OM$, and $\\OL$. In all plots the unit of distance is taken to be $c/H_0$. In these figures $D_{\\ell}$ is plotted on a magnitude scale, 5 Log $D_{\\ell}$ (\\ie the distance modulus plus 5\\,Log $10pc\\, H_0/c$). In Figure 4, $\\OL$ is held fixed while $\\nu$ and $\\OM$ are varied and in Figure 5, $\\OM$ is held fixed while $\\nu$ and $\\OL$ are varied. In Figure 6, $\\OO=\\OM+\\OL=1$ is fixed while all three parameters vary. In Figure 7 the sensitivity of observed magnitudes to variations of $\\OM$ is illustrated by fixing $z=0.83$ and $\\OL=0.1$. In Figure 8 a similar plot is given showing the sensitivity to variations of $\\OL$. The importance of the clumping parameter is easily seen from these last two figures. If the distance modulus of a source such as SN 1997ap at $z=0.83$ were precisely known (\\eg see the two sample horizontal lines in Figure 7) then a determination of $\\OM$ could be made, assuming $\\OL$ were somehow known. Likewise, from Figure 8, a determination of $\\OL$ could be made if $\\OM$ were somehow known. From Figure 7 the reader can easily see that the determined value of $\\OM$ depends on the clumping parameter $\\nu$. The $\\OM$ value will be about 95\\% larger for a $\\nu=2$ completely clumpy universe than it will be for a $\\nu=0$ completely smooth FLRW universe. Equivalently, $\\OM$ could be underestimated by as much as 50\\% if the FLRW is used. The maximum underestimate is reduced to 33\\% at the smaller redshift of $z=0.5$ (see a similar result for $\\OL=0$ in \\cite{KVB}). These conclusions are not sensitive to the value of $\\OL$. Slightly different conclusions follow from Figure 8 about $\\OL$. The discrepancy in the determined value of $\\OL$ is $\\Delta \\OL \\sim -0.14$ for $\\nu=2$ compared to $\\nu=0$, and is not sensitive to the distance modulus. The discrepancy is halved, $\\Delta \\OL \\sim -0.07$, at a smaller redshift of $z=0.5$. A minimal estimate of the quantity of data required to begin distinguishing between the various $\\nu$ values can easily be made. At $z=0.5$ the differences in observed magnitudes of a SN Ia in a $\\nu=0$ (100\\% smooth FLRW) and a $\\nu=2$ (100\\% clumpy) universe is about $\\Delta m \\sim 0.02$ if $\\OM\\sim 0.2$, and $\\Delta m \\sim 0.09$ if $\\OM\\sim 0.8$. These differences are not sensitive to $\\Lambda$. With corrected-intrinsic and observed magnitude uncertainties of $\\pm 0.2$,\\ \\cite{BD}, data on over 200 SNe Ia will be required if we live in a low density universe and over a dozen if we live in a higher density one. The results presented here (\\ref{Dell0}),(\\ref{Dellinfty}),(\\ref{Dell01}), and (\\ref{Dellinfty1}) for the `intergalactic' distance--redshift relation are quite general. They contain corrections (for mass inhomogeneities) to the standard FLRW result, applicable to observations where gravitational lensing can be neglected, \\ie observations where the conformal (Weyl) curvature doesn't produce significant average shear in (\\ref{Af}). Even though the original area equation (\\ref{Af}) was rigorously established for a particular type of Swiss cheese model, the resulting equation which neglects lensing (\\ref{Ab0}) is expected to be widely applicable to observations at redshifts of $z=1$ and less. Application of its solution to a given set of observations requires that the average fraction of the mass density contained in the observing beams (\\ie the $\\nu$ parameter) be determined. This fraction obviously depends on the number as well as the type of object observed. Collecting CMB radiation at wide angles is likely to produce a $\\nu=0$ value but observing a few dozen SNe Ia might well result in a value close to $\\nu=2$ (\\ie we might in fact live in a universe where mass, dark or otherwise, is primarily associated with galaxies). If a significant fraction of the universe's mass density is clumped on galactic scales, then the effects of these clumps on SNe Ia observations should be taken into account by using the lensing formulas rather than by decreasing $\\nu$ to zero. Recent numerical work by \\cite{HD} confirms the assertion that, given galaxy clumping, the cross-sectional area of a \\underbar{typical} light beam will not follow the FLRW area-redshift relation. Instead the area will follow more closely one of the `intergalactic' \\mz relations given here, until a lensing event occurs. The new luminosity distances presented here represent the theoretical minimum of the observed magnitudes and are especially applicable to situations where lensing is infrequent (\\ie where the most probable value is closer to the min than the mean). Because \\cite{HD} did not include any diffuse transparent matter, the applicable \\mz relations given here are those with $\\nu=2$. For $\\Lambda=0$ and $\\Omega_0<1$ it is the Dyer-Roeder solution (\\ref{DC1}) and for $\\Lambda=0$ and $\\Omega_0=1$ its the $\\nu=2$ solution of Dashevskii \\& Slysh (\\ref{Dash}). The $\\nu=0$ (standard FLRW) result represents the theoretical `mean' for \\mz for a universe in which only weak-lensing events occur. For extremely non-symmetric probability distributions, the ``mean\" is not likely the best estimator - in this case the ``most probable\" is likely better, \\cite{SD}." + }, + "9802/astro-ph9802178_arXiv.txt": { + "abstract": "We present optical and infrared observations of \\rp, a faint ($I = 21.01$) and very red object ($I - K = 4.57$) discovered in a deep CCD survey, covering an area of 800 square arcmin of the Praesepe open cluster. A low resolution spectrum shows that \\rp\\ is a very late object, the latest object in Praesepe for which a spectrum has been taken to date. Our estimates give a mass between 0.063 and 0.084 $M_{\\sun}$, and indicate that \\rp\\ may turn out to be the first brown dwarf in this cluster. ", + "introduction": " ", + "conclusions": "" + }, + "9802/astro-ph9802334_arXiv.txt": { + "abstract": "We describe a method for determining the limb polarization and limb darkening of stars in eclipsing binary systems, by inverting photometric and polarimetric light curves. Because of the ill-conditioning of the problem, we use the \\relax Ba\\-ckus-Gil\\-bert\\relax {} method to control the resolution and stability of the recovered solution, and to make quantitative estimates of the maximum accuracy possible. Using this method we confirm that the limb polarization can indeed be recovered, and demonstrate this with simulated data, thus determining the level of observational accuracy required to achieve a given accuracy of reconstruction. This allows us to set out an optimal observational strategy, and to critcally assess the claimed detection of limb polarization in the Algol system. The use of polarization in stars has been proposed as a diagnostic tool in microlensing surveys by Simmons et al.\\ (\\cite{simmons95}), and we discuss the extension of this work to the case of microlensing of extended sources. ", + "introduction": "Scattered light emerging from a stellar atmosphere is expected to be partially linearly polarized. This effect should be greatest at the limb of the star, with a pure electron scattering atmosphere giving a limb polarization of 11.7\\% (Chandrasekhar \\cite{chandra46}, \\cite{chandra50}). Less idealised calculations (Collins \\& Buergher \\cite{Collins+Buerger}) suggest a lower degree of limb polarization, around 2\\% for early-type stars, though this figure is rather sensitive to the ionization state of the outermost atmospheric layers. Clearly a spherically symmetric star will exhibit no net polarization, but if the symmetry is broken by an eclipse the limb polarization should in principle be detectable. Theoretical polarimetric light curves for this situation have been calculated by Landi Degl'Innocenti et al.\\ (\\cite{landi88}). The first detection of this `Chandrasekhar effect' was in the \\object{Algol} system (Kemp et al.\\ \\cite{kemp83}). This data was analysed by Wilson and Liou (\\cite{wilson93}), but the complexity of the system and the amount of modelling involved in the analysis prevented them from making any reliable estimate of the limb polarization. The inversion of polarimetric light curves from eclipsing binary stars should allow the limb polarization of the eclipsed star to be measured (and the photometric light curve should similarly give the limb darkening). In fact this inverse problem is highly ill-conditioned, and relating observations to stellar atmosphere models is therefore far from straightforward. In this paper we investigate the practical feasibility of determining of limb polarization by this method. This allows us to address three closely related issues: \\begin{enumerate} \\item We develop a method of obtaining the polarization at a point on the stellar disc, and of estimating the error on this value. This is based on the \\relax Ba\\-ckus-Gil\\-bert\\relax {} inversion technique. \\item We determine the maximium accuracy possible in determining limb polarization, given a number of data points and a noise level. \\item Thus, we are able to put forward an observational strategy which should allow the best measurement of limb polarization. \\end{enumerate} In section 2 we give a brief overview of the problem. We set out the formalism of the \\relax Ba\\-ckus-Gil\\-bert\\relax {} method in section 3, and discuss its suitability for the problem at hand. Section 4 contains the calculations for the specific case of eclipsing binary stars, and section 5 presents the results of the inversion scheme when applied to simulated data, and the conclusions that can be drawn from these. Section 6 considers a simplified analogue of the Algol system, comparing the theoretical polarization profile with the best resolution current measurements can achieve. ", + "conclusions": "Our studies of simulated data show the fundamental limitations on the determination of limb polarization in eclipsing binary stars. In particular, \\relax Fig.~\\ref{f:lambdas} indicates that limb polarizations of order a few percent are only just above the threshold of detectability, even in perfectly spherically symmetric, non-interacting binary systems. The situation will be worse in more complex systems. It is important to appreciate that the \\relax Ba\\-ckus-Gil\\-bert\\relax {} method does not strictly estimate the polarization at a point on the stellar disc, but rather the polarization convolved with the resolution function. To relate the results of the inversion to a particular model, it is necessary to calculate the theoretical value of this convolution, which should be consistent with the $\\lambda$ we have chosen and with all higher values of $\\lambda$ -- these represent coarser averages over the stellar disc. The bottom line is that one must take care in drawing conclusions about limb polarization from studies of eclipsing binaries. It is clearly not possible to distinguish between stellar atmosphere models on this basis if their predicted limb polarizations differ by less than the maximum accuracy achievable. On the other hand, an appreciation of the issues raised in this paper will allow a meaningful determination of limb polarization, with reliable error estimates. Formally, the eclipsing binary problem is very similar to the gravitational microlensing problem. Indeed, part of the initial motivation for this work sprang from studies of the microlensing of extended sources. A future paper (Coleman et al.\\ \\cite{CGS2}) will apply the inverse problem approach outlined in this paper to the use of microlensing as a probe of stellar atmospheres." + }, + "9802/astro-ph9802102_arXiv.txt": { + "abstract": "\\noindent We discuss the spin evolution of pulsars in the case where a superfluid component of the star is coupled to the observable crust on long, spindown timescales. The momentum transfer from the superfluid interior results in an apparent decay of the external torque and, after a dramatic increase, to an asymptotic decrease of the generic value of the braking index, e.g. $n=3$, to values $n\\sim 2.5$ if the magnetic field of the star does not decay over its lifetime. In the case where an exponential decay of the magnetic field towards a residual value occurs, the star undergoes a spin-up phase after which it could emerge in the millisecond sector of the $P$-$\\dot P$ diagram. ", + "introduction": " ", + "conclusions": "" + }, + "9802/astro-ph9802272_arXiv.txt": { + "abstract": "This paper presents 3600---5400\\,\\AA\\ integrated spectra of 19 \\gcs\\ (or candidates) projected on the central regions of M\\,31, $r\\leq5.3$\\min\\ ($\\approx$1.2\\,kpc). We check the cluster nature of these objects, and derive their ages, metallicities and reddenings. From the initial sample, 16 objects turn out to be true star clusters, two are Galactic dwarf stars, and one is a high redshift background galaxy. Only two clusters are found to be super \\mr, suggesting that this phenomenon is not very common. For some clusters, we cannot rule out the possibility that they are of intermediate age; this requires confirmation by observations at the calcium triplet. We also present the metallicity histogram of this central bulge sample and discuss possible scenarios to explain its properties. ", + "introduction": "Globular clusters in the Galaxy and in the Local Group play a fundamental role in connecting studies of individual stars and integrated properties of star clusters. In this respect, studies of star clusters in the Andromeda galaxy (M\\,31) are particularly suitable. The proximity and inclination of this giant spiral allow one to study in detail each of its constituent sub-systems, halo, disk, and bulge even near the galaxy center. M\\,31 has been already extensively surveyed for \\gcs\\ e.g., Vetesnik (1962), Sargent et al. (1977) and the Bologna group catalogue by Battistini et al. (1987). The brightest clusters have been explored by means of integrated spectroscopy, which makes it possible to determine some of their fundamental parameters, such as metallicity, age and reddening, as well as to confirm their true nature as \\gcs\\ (e.g. Burstein et al. 1984, Tripicco 1989, Huchra, Brodie \\& Kent 1991, Jablonka, Alloin \\& Bica 1992, Bica et al. 1992, and Huchra et al. 1996). A number of studies have concentrated on identifying clusters near the center of M31, e.g., Alloin, Pelat \\& Bijaoui (1976), Auri\\`ere, Coupinot \\& Hecquet (1992) and Battistini et al. (1993). In the present paper, we concentrate on the central 10\\min\\ by 10\\min\\ region of M\\,31 and derive velocities, metallicities, ages and reddenings of the confirmed star clusters. This work is a first attempt at getting a picture of the properties of \\gcs\\ in the inner bulge of a spiral galaxy (other than our own), which in turn must help constrain models of galaxy formation and evolution. One of our principal goals is to determine how frequent the occurrence of super \\mr\\ \\gcs\\ can be. Bulge central regions, locations of intense star formation, appear to be one of the best places to find them. The methods employed in the present work, i.e. comparison with template cluster spectra and grids of spectral properties as a function of age and metallicity, have been developed and extensively used to study star clusters in our Galaxy (e.g. Bica \\& Alloin 1986a,b), the Magellanic Clouds (e.g. Santos et al. 1995a), M\\,31 itself (Jablonka et al. 1992), and more recently NGC\\,5128 (Jablonka et al. 1996). This paper is organised as follows: Section\\,2 gives a description of the sample; Section\\,3 details observations and reductions; Section\\,4 deals with the radial velocity measurements; Section\\,5 presents measurements of equivalent widths and results for cluster parameters; a discussion of the results is given in Section~6. Finally, we summarise our conclusions in Section\\,7. ", + "conclusions": "We have carried out spectroscopy of a sample of 19 objects projected on the bulge of M\\,31: 16 were confirmed as star clusters, 2 are Galactic dwarf stars, and 1 is a high redshift background galaxy. The objects confirmed as star clusters had been classified as Bologna classes A and B, whereas the 3 non-clusters were of the low probability class C. We derived radial velocities, metallicities, ages and reddenings for our sample clusters. In some cases, especially G\\,175, there remains an ambiguity between intermediate and old cluster ages, due to the limited wavelength range observed. Two clusters are found to be super \\mr, G\\,174 and G\\,177. Their features are comparable to those in the most \\mr\\ stellar populations, i.e., nuclei of giant ellipticals and the semi-stellar nucleus of M\\,31 itself. The metallicity distribution of the sample is compatible with the distribution of globulars in the bulge of our Galaxy. Although small, both the M\\,31 and the Galaxy \\gc\\ samples seem to have an excess of \\mp\\ objects with respect to the stars in the Galactic bulge. Some interloping halo \\gcs\\ might be present in both cluster samples. Alternatively, this might be due to different formation histories. In the future, it will be important to observe clusters surrounding the nucleus in the other quadrants, in order to increase the statistical significance of the properties analyzed. In particular this would make it possible to check for the existence of rotation in the cluster system, as well as to identify new super \\mr\\ clusters, imposing constraints on chemical evolution models, especially at the high metallicity end." + }, + "9802/astro-ph9802044_arXiv.txt": { + "abstract": " ", + "introduction": "The small but growing subpopulation of $\\gamma$-ray active radio pulsars raises our hopes in advancing the theoretical models of pulsar emission. In this paper we concentrate on the picture in which both the particle acceleration and $\\gamma$-ray emission occur in the inner magnetosphere of a neutron star (NS) above the magnetic pole. The corresponding models, known as polar cap (PC) models, favor relatively small obliquities (single pole emission), and they crucially depend on curvature radiation (CR) and on the QED process of magnetic pair creation. The prototype of the PC models has been suggested by Ruderman \\& Sutherland (1975). Some outstanding problems and progress in this field were addressed clearly by Curtis Michel, Jonathan Arons, and other speakers at the conference. In this written report we will schematically describe the relevant results of our most recent work that are being prepared for publication in a more complete form. The bottom line of this presentation is that {\\bf the PC model for pulsar high-energy emission is generally a viable approach, but it needs to be revised and improved in a number of important theoretical aspects.} ", + "conclusions": "In conclusion we emphasize that the PC model is capable of explaining and reproducing (see Harding 1996, Daugherty \\& Harding 1996) the main observational facts mentioned in $\\S~2$: the widely separated double-peaked profiles (also in hard X-rays) with bridge component, very steep high-energy spectral cutoffs (due to magnetic pair production), and systematic soft-hard-soft hardness variation in the pulse (due to the softening of emerging spectra by cascades). The novel developments we would like to introduce here are: 1) the double PFF controlled by CR, a self-limiting buildup where the electrons are accelerated from the lower front and a small fraction of positrons returning from the upper front to produce the lower front; 2) the establishment of the PFFs and pair cascades at higher altitudes; and 3) the revised treatment of a feedback between the inflow of cascade particles into the acceleration region and accelerating electric field, including the PC heating by returning positrons. Although the PC model has had success in accounting for a number of main observational features of $\\gamma$-ray pulsars, there remain some puzzling observational data and exceptions (e.g., Crab: alignment of $\\gamma$-ray pulses with those at other wavebands; Geminga: complex behaviour of spectral hardness through the pulse; Geminga, Vela, and PSR 0656+14: soft pulsed X-ray emission; etc.). Given the advancement in understanding of PC acceleration discussed above, we will be able to model characteristics of pulsar high-energy emission from PC pair cascades in more detail and ultimately address these unsolved problems." + }, + "9802/astro-ph9802116_arXiv.txt": { + "abstract": "{We present new data on the photodissociation regions associated with the reflection nebula NGC~7023, particularly the three bright rims to the north, south and east of the illuminating star HD 200775. $^{13}$CO(3--2) emission, mapped at 20\\arcsec\\ resolution at the Caltech Submillimeter Observatory (CSO), delineates a molecular cloud containing a cavity largely devoid of molecular gas around this star. Neutral carbon is closely associated with the $^{13}$CO emission while ionized carbon is found inside and at the edges of the cavity. The ionized carbon appears to be, at least in part, associated with \\ion{H}{1}. We have mapped the northern and southern rims in $^{12}$CO(6--5) emission and found a good association with the H$_2$ rovibrational emission, though the warm CO gas permeates a larger fraction of the molecular cloud than the vibrationally excited H$_2$. The column density contrast between the bright rims and the diffuse region inside and in front of the cavity is about 10. Despite the fact that the edges of the cavity are viewed edge-on, the carbon emission extends much further into the molecular gas than does the photodissociation region, as defined by the H$_2$ emission region. Geometrically, NGC 7023 consists of a sheet of dense molecular gas in which the star was born, subsequently blowing away much of the surrounding gas. The three bright rims are located at the edges of the remaining molecular cloud, and are viewed approximately edge-on. The results are compared with PDR models, invoking direct illumination from the star, which are largely successful, except in explaining the presence of neutral carbon deep in the molecular cloud. We suggest that, in the particular case of NGC~7023, a second PDR has been created at the surface of the molecular cloud by the scattered radiation from HD 200775. This second PDR produces a layer of atomic carbon at the surface of the sheet, which increases the predicted [C]/[CO] abundance ratio to 10\\%, close to the observed value. Further tests for the applicability of PDR models in such regions are suggested. } ", + "introduction": "Reflection nebulae are produced when a massive star illuminates a molecular cloud. Due to the enhanced radiation field, the gas surrounding the star is heated and its chemistry is modified. Reflection nebulae belong to the generic class of sources known as PDRs (Photon-Dominated Regions). These are typically located at the transition layer between warm, ionized gas produced by an intense radiation field and a cold neutral atomic or molecular cloud. If this definition is extended to low values of the radiation fields, down to the average intensity of the interstellar radiation field (ISRF), $I = G_0 \\times \\rm ISRF$ with $G_0 = 1$ at the solar radius, PDRs comprise a very significant fraction of the mass of the neutral gas in our Galaxy (Hollenbach \\& Tielens 1995). Furthermore, due to their high efficiency in producing intense emission (line and continuum), PDRs produce a significant fraction of the dust and gas emission of external galaxies including the CO rotational lines and \\ion{C}{1} \\& \\ion{C}{2} fine structure lines (Bennett et al.\\ 1994, Hollenbach \\& Tielens 1995). In PDRs at the surface of molecular clouds illuminated by a strong radiation field, $G_0$ $\\simeq$ 1000 or more, the gas temperature is high and, in general, atomic and molecular lines are intense. For some reflection nebulae, the geometry of the cloud is visually revealed. Such sources are thus suitable places to study basic physical processes of the interstellar medium. In particular the processes leading to the thermal balance of interstellar gas can be investigated in detail due to the knowledge of the radiation field and the overall geometry of the source. Carbon plays a key role in PDR models because it is expected to be dominantly in three different species throughout the PDR\\@. There is a layer of almost totally ionized carbon at the outer edge, an intermediate region where carbon is neutral and then the cold molecular interior where carbon should be locked into CO\\@. Because oxygen and nitrogen remain neutral in the outer edge, no equivalent structure can be seen in the main repositories of these other abundant species so carbon plays a unique role in testing PDR models. The possibility of observations of C$^+$, C and CO at the same spatial locations in the same source therefore provides a unique diagnostic of predictions: the depth and width of the CI emission zone can test the accuracy of the chemistry and radiative transfer, and the excitation of C and CO are sensitive to the thermal balance. Previous observations of interstellar clouds reveal that atomic carbon is generally both overabundant and more widely distributed than simple models predict (Phillips \\& Huggins 1981, Keene et al.\\ 1985). This has led to questions concerning the homogeneity of the cloud resulting in proposals for porous clouds with greater than expected UV penetration (Phillips \\& Huggins 1981, Stutzki et al.\\ 1988). In that case a PDR structure is expected at the surface of each substructure in the porous cloud as long as a significant UV field is present ($G_0 \\geq 1$ to produce \\ion{C}{1} (Spaans 1996)). A widely used test for models of the surface of the PDR is the H$_2$ rovibrational emission in the near infrared. H$_2$ reveals details of the external zone of the PDR, but because the emission is very sensitive to both the UV field intensity and the gas density, the extent of the H$_2$ emission is much reduced compared to that of the three carbon species, even in the case of a porous cloud. The combination of data from both the H$_2$ rovibrational emission and the carbon budget therefore represents a powerful tool for investigating PDRs, particularly at high spectral resolution when the velocity structure can be resolved. We present in this paper a new $^{13}$CO (3--2) map of the reflection nebula NGC~7023 together with selected data of C$^{18}$O(3--2), $^{12}$CO(6--5), neutral (\\ion{C}{1}) and ionized (\\ion{C}{1}) carbon emission. A $^{12}$CO(3--2) map has been taken but is not shown, because the structure of the cavity region is obscured by the lower column density foreground material. Though the NGC~7023 reflection nebula has been extensively investigated in the past, this is the first observation of neutral carbon and the first study of \\ion{C}{2} at high spectral resolution. NGC 7023 is illuminated by the young B3Ve star HD 200775 ($\\alpha(1950) = 21^{\\rm h}00^{\\rm m}59\\fs 7, \\delta(1950) = 67\\arcdeg 57\\arcmin 55\\farcs 5$), at a distance of 600 pc (Rogers, Heyer, \\& Dewdney 1995). The radiation field is enhanced by a factor of a few thousand over the average value in the solar vicinity (Chokshi et al.\\ 1988). Due to this enhanced radiation field, the very small dust grains experience large excursions in temperature and produce continuum emission in the near IR, as first observed by Sellgren (1984) for this and other reflection nebulae. The scattering properties of dust grains have also been investigated in NGC 7023 (Murthy et al.\\ 1993, Witt et al.\\ 1993). Maps of the dust continuum emission in the far infrared (Whitcomb et al.\\ 1981, Casey 1991) have been taken at a spatial resolution of about 1\\arcmin. The dust temperature varies from 50 K in the vicinity of HD 200775 to 20K in the molecular cloud. The decrease in dust temperature is accompanied by an increase of the dust opacity at 250 $\\mu$m, indicating a rise of column density in the molecular cloud surrounding HD 200775. The structure of the molecular gas has been determined at low spatial resolution (1\\arcmin\\ -- 2\\arcmin) by Watt et al.\\ (1984), Fuente et al.\\ (1992) and Rogers et al.\\ (1995) by mapping in $^{12}$CO, $^{13}$CO and C$^{18}$O. The $^{13}$CO emission is found at the border of the bright nebula, and delineates a cavity with an hourglass shape roughly centered on HD 200775, in which the $^{13}$CO and C$^{18}$O emission is extremely weak or not detectable. This is seen in Figure 1, which is our new high spatial resolution ($\\sim$ 20\\arcsec) $^{13}$CO(3--2) map of most of the region. At this resolution the bright inner edge of the $^{13}$CO emission traces the nebulosity seen on the POSS plate. Rogers et al.\\ (1995) were also able to isolate the \\ion{H}{1} emission from the nebula from the general, widespread \\ion{H}{1} emission. They found \\ion{H}{1} emission from the cavity, mostly north of the star, and conclude that there is global pressure equilibrium between the warm gas around the star and the gas in the molecular cloud. \\ion{H}{1} is slightly redshifted compared to the molecular gas, by about 2 km$\\;$s$^{-1}$. Fuente et al.\\ (1996) obtained an \\ion{H}{1} map with 10\\arcsec\\ resolution which showed that the neutral gas does not completely fill the interior of the cavity, but accumulates on the inner edges, at the surface of the molecular gas, as expected for a stratified PDR\\@. The absence of significant \\ion{H}{1} emission in the largest lobe of the cavity, west of the star, and the presence of background stars indicates that we are dealing with a sheet of material in which the star HD 200775 has created a nearly complete hole (Rogers et al.\\ 1995). Fuente et al.\\ (1993) have investigated, at high spatial resolution ($\\sim$ 12\\arcsec), the northern part of the nebula, 1\\arcmin\\ north-west of the star. They found a sharp ridge in $^{13}$CO and C$^{18}$O emission, which we here call the north rim. They have shown that the gas chemistry is affected by the UV radiation and that abundance ratios such as [CN]/[HCN] and [HCN]/[HNC] are enhanced in the gas affected by the UV\\@. As deduced from HCN measurements, the density in the gas is fairly high, about 10$^5$ cm$^{-3}$ in the ridge. Somewhat lower densities ($0.3 - 2 \\times 10^4$ cm$^{-3}$) are obtained from the excitation of CO and its isotopes. Ionized carbon and atomic oxygen have been detected at low spectral resolution by Chokshi et al.\\ (1988) in the nebula, and they presented the first PDR model for NGC 7023. The PDR model has been refined by Lemaire et al.\\ (1996), who obtained images of the rovibrational lines of molecular hydrogen with 1\\arcsec\\ resolution. The H$_2$ emission is concentrated in narrow and long filaments, located mostly north of the star. These filaments have also been detected in HCO$^+$ (Fuente et al.\\ 1996). Lemaire et al.\\ (1996) show that the H$_2$ emission is not uniform in the filaments but shows variations down to the 1\\arcsec\\ scale. Both the high brightness and spatial structure of the H$_2$ emission are evidence that the illuminated gas is fairly dense, $n_{\\rm H} \\simeq 10^5~ {\\rm cm}^{-3}$ (Lemaire et al.\\ 1996, Martini, Sellgren, \\& Hora 1997). There is a second H$_2$ front south of the star, fainter than the northern front by a factor 2--3. Similar filaments are found in R and V-R pictures (Watkin, Gledhill, \\& Scarott 1991) as sources of Extended Red Emission (ERE). H$_2$ emission and ERE are coincident in both fronts (50\\arcsec\\ NW and 70\\arcsec\\ S) but have different small scale structure (Lemaire at al.\\ 1996). It is clear from Figure 1 that a third powerful PDR exists on the east rim of the cavity. We have initiated some studies here, but in total there is little data available on this. It should be investigated further in the future. Because high $J$ CO lines and the carbon fine structure lines are predicted to be strong by these PDR models, we have observed the $J$ = 3--2 lines of $^{12}$CO, $^{13}$CO and C$^{18}$O, the $J$ = 6--5 line of $^{12}$CO and the 1--0 line of \\ion{C}{1} in NGC 7023 with the Caltech Submillimeter Observatory (CSO). We also present Kuiper Airborne Observatory (KAO) spectra of the \\ion{C}{2} emission at high spectral resolution. We describe the observations in the next section, then present the results and compare the line maps with the H$_2$ images in order to obtain a more precise understanding of the geometry of the nebula. Finally we discuss the validity of the PDR models for this source. ", + "conclusions": "PDR models are well established as capable of explaining the molecular and atomic emission lines seen on the boundaries of molecular clouds when illuminated by starlight (Hollenbach \\& Tielens 1995). These models have recently been modified to incorporate new data on reaction rates and molecular excitation processes. We have used the PDR model based on the code of Abgrall et al.\\ (1992) and Le Bourlot, Pineau des For\\^ets, \\& Roueff (1993a), with different geometric configurations, in order to check the validity of actual parameters and to investigate the marked difference between the northern and the southern rims in NGC 7023, in both the H$_2$ and \\ion{C}{1} intensities relative to $^{13}$CO\\@. In part these differences could be due to geometrical effects. Since this PDR model starts with a semi-infinite H$_2$ slab on the surface of which the UV impinges, it cannot address the structure of the cavity, the region between the PDR and the star. >From Figure 5, it is clear that the north rim is viewed almost edge-on: the large variation of the column density of molecular hydrogen as traced by $^{13}$CO(3--2) emission, the position of the molecular hydrogen emission at the edge of the molecular cloud, and the coincidence of the peak in the \\ion{C}{1}/$^{13}$CO ratio with the zone of H$_2$ emission result naturally from the edge-on geometry. It is quite likely that the actual geometry is that of a sheet with a sharp cut-off edge, which is the surface illuminated by the star. The sections of these edges closest to HD 200775 are the north and south rims, where bright H$_2$ emission has been detected. In order to investigate the role of the viewing geometry in a PDR model, we have calculated abundances as a function of depth within a bright rim. These abundances are naturally interpreted as corresponding to a ``face-on'' viewing geometry, and from them we approximated the emergent intensities for a PDR viewed ``edge-on'' by summing the local emissivities at each depth inside the cloud along a path length of $3 \\times 10^{16}$ cm (0.01 pc). This approximation is valid only for optically thin lines for which the emissivities may be enhanced significantly by limb brightening. The model we have used is an isochoric model, with a UV enhancement factor $G_0$ of 10$^4$ and a density of $n{\\rm (H + 2H_2) = 2 \\times 10^5~cm^{-3}}$. We have used as gas phase elemental abundances the values found towards $\\zeta$ Oph listed in Table 2. We have also tried a model with oxygen, carbon and nitrogen depleted by a factor of three from these values. From Lemaire et al.\\ (1996), it is known that the hydrogen density has to be larger than about 10$^5$ cm$^{-3}$ to reproduce the observed intensities of the H$_2$ rovibrational lines. An analysis of the molecular emission lines (CS and HCN) also suggest high densities, in the 10$^5$ cm$^{-3}$ range (Fuente et al.\\ 1993). The results are illustrated in Figure 7 where we show the variation of the local abundances $n({\\rm X})/(n({\\rm H}) + 2n({\\rm H_2}))$ of $^{12}$CO, $^{13}$CO, C and C$^+$ with the depth into the cloud expressed in arcsec at the distance of NGC 7023 (600 pc) (top panel) and the predicted emissivities for the edge-on view (bottom panel). For the optically thick $^{12}$CO(3--2) line, little or no limb brightening is expected and we show only the face-on emissivity. In this plane-parallel structure, the model predicts a layered structure with C$^+$ at the outer boundary of the cloud, then C and CO\\@. For this isochoric model, in the outer region where C and C$^+$ dominate over $^{12}$CO, the high density and temperature cause efficient excitation of $^{12}$CO and therefore produces strong high $J$ $^{12}$CO emission. We discuss below the different layers in the photo dissociation region: 1) \\ion{C}{2} emission is found in the warm ionized and partially dissociated medium which lies at the surface on which the UV impinges. It is also present in the next layer where H$_2$ is able to resist photodissociation and is strongly excited. This is qualitatively consistent with the observed C$^+$ data which show two different components depending on the actual observed positions. The emission from the \\ion{C}{2} region in the models would be clearly diluted in the 45\\arcsec\\ KAO beam. With a dilution factor of about 5, i.e.\\ assuming that the \\ion{C}{2} emission comes from a structure of width 10\\arcsec\\ and longer than 60\\arcsec, the observed peak antenna temperature of the \\ion{C}{2} emission from this region corresponds to an absolute antenna temperature of 75 K, close to the predicted values. 2) The next layer is the atomic carbon region and the main problem for the models is the fit of the carbon data: whereas the model predicts no \\ion{C}{1} emission at distances larger than 20\\arcsec\\ from the PDR, \\ion{C}{1} has been observed throughout the molecular cloud (Fig. 5). The distribution of neutral carbon is therefore much more extended than expected for a simple PDR\\@. Because reflected starlight is seen in this nebula, and neutral carbon is produced in even low radiation field environments, it is possible that some neutral carbon emission we observe comes from the externally illuminated visible surface of the sheet. We have evaluated the UV intensity due to scattered light from HD 200775 using the observations with UIT (Witt et al.\\ 1992). We found an enhancement factor of $G_0 \\sim 20$ at 100\\arcsec\\ from HD 200775. This scattered radiation forms another PDR at the surface of the molecular cloud. To evaluate the role of this second PDR in producing neutral and ionized carbon at large distances from the north rim, we have run a PDR model with the same abundances, the same high density, $n = 2 \\times 10^5$ cm$^{-3}$, but a low radiation field $G_0$ = 50. The variation of the abundances of C$^+$, C and CO as a function of the depth into the cloud A$_V$ are displayed in Figure 8. For A$_V$ of a few magnitudes, this second PDR is an important source of neutral carbon in the molecular gas, so that the predicted [C]/[CO] abundance ratio is larger than 10\\%. Indeed, even in a low UV field, the PDR is associated with ionized and neutral carbon layers at the surface of the molecular cloud. When viewed ``face-on'' the layered structure results in significant cumulative column densities of neutral and ionized carbon for all values of A$_V$ hence significant abundances of C$^+$ and C appear, even though most of the neutral and ionized carbon reside at the outer surface. With this low UV flux, the model does not predict any H$_2$ or warm dust emission, as observed. In further studies this proposal could be tested by determining if a substantial amount of C$^+$ is also present in the mixed C/CO regime, which could indicate a UV origin to the effect. This would require a platform such as SOFIA or FIRST with heterodyne receivers. PDR models predict the existence of a zone where carbon isotopic fractionation takes place at the edge of the molecular cloud close to the UV source. In that zone, the [$^{12}$CO]/[$^{13}$CO] abundance ratio drops to 30 (Kopp et al.\\ 1996). Carbon fractionation takes place as soon as the formation of molecules enhances the cooling and the temperature drops. It stops deeper into the cloud because of the lack of C$^+$. If the PDR model is appropriate for these regions far from the star, $^{13}$CO/$^{12}$CO fractionation should exist in such a regime and could be observed with ground based techniques. If carbon deeper into the cloud is due to non-UV processes, such as turbulent mixing (Xie, Allen, \\& Langer 1995) or high ionization phase chemistry (Le Bourlot et al.\\ 1993b; Le Bourlot, Pineau des For\\^ets, \\& Roueff 1995; Flower et al.\\ 1994), there will be much less C$^+$ and fractionation. The assumed values for elemental abundances are important parameters for PDR models. The models with depleted elements exhibit the same brightness for the $^{12}$CO and $^{13}$CO lines essentially due to the high optical depths, but reduced emissivities of ionized and neutral carbon by a factor roughly equal to the input depletion. The strength of both the observed \\ion{C}{1} and \\ion{C}{2} lines favor the models with rather high gas phase abundances of carbon and oxygen, i.e.\\ the models using the $\\zeta$ Oph abundances. Because the cooling is slightly reduced in the models with lowered elemental abundances, the $^{13}$CO edge is not as sharp as in the case of $\\zeta$ Oph abundances. In any case, the edge is likely to be unresolved because the emission rises over $3 \\times 10^{16}$ cm. At the distance of NGC 7023, 450 -- 600 pc, this corresponds to at most 4\\arcsec, which is below the resolution of these observations but could be resolved by millimeter interferometers." + }, + "9802/hep-ph9802423_arXiv.txt": { + "abstract": "\\widetext We study random bubble lattices which can be produced by processes such as first order phase transitions, and derive characteristics that are important for understanding the percolation of distinct varieties of bubbles. The results are relevant to the formation of topological defects as they show that infinite domain walls and strings will be produced during appropriate first order transitions, and that the most suitable regular lattice to study defect formation in three dimensions is a face centered cubic lattice. Another application of our work is to the distribution of voids in the large-scale structure of the universe. We argue that the present universe is more akin to a system undergoing a first-order phase transition than to one that is crystallizing, as is implicit in the Voronoi foam description. Based on the picture of a bubbly universe, we predict a mean coordination number for the voids of 13.4. The mean coordination number may also be used as a tool to distinguish between different scenarios for structure formation. ", + "introduction": " ", + "conclusions": "" + }, + "9802/astro-ph9802283_arXiv.txt": { + "abstract": "We investigate the effect of inverse-Compton scattering of flares of soft radiation in different geometries of a hot, Comptonizing region and a colder accretion disk around a solar-mass black hole. The photon-energy dependent light curves, their Fourier transforms, power spectra and Fourier-period dependent time lags of hard photons with respect to softer photons are discussed. On the basis of a comparison with existing data we find arguments against Comptonization of external soft radiation as well as Comptonization in a homogeneous medium as dominant mechanisms for the rapid aperiodic variability in Galactic black-hole candidates. Possible further observational tests for the influence of Comptonization on the rapid aperiodic variability of Galactic black-hole candidates are suggested. ", + "introduction": "The rapid aperiodic and quasiperiodic variability of the X-ray signals from Galactic black-hole candidates (GBHCs) and low-mass X-ray binaries (for a review see \\cite{vdk95}) contains valuable information about the source of high-energy emission in these objects. Its typical timescales, typical repetition frequency, power spectrum density (PSD), autocorrelation function, time lags between different energy bands etc. give hints toward important parameters such as the extent of the emitting region, the dominant microscopic timescales, the dominant emission mechanism and the geometry of the source. With the new generation of X-ray telescopes such as the PCA and the HEXTE on board RXTE it is now possible to measure the above properties of the rapid variability of X-ray binaries in great detail with high timing resolution and good spectral resolution. As we will see in this paper, the photon-energy dependence of the rapid aperiodic variability can provide important diagnostics of the nature of the X-ray emitting regions. Early measurements of the Fourier-frequency dependence of time lags between the signals in different X-ray energy channels from GBHCs (\\cite{mkk88}, \\cite{mik93}) have been interpreted as evidence that Comptonization of soft photons in a hot, uniform plasma could not be the dominating mechanism for the production of hard X-rays. This conslusion, however, is strongly geometry-dependent and does not hold for a very extended, inhomogeneous Comptonizing region, as was recently found by Kazanas et al. (1997). In that and in two subsequent papers (\\cite{hkt97a}, \\cite{hkc97b}) the effect of radial density gradients in the hot Comptonizing regions in GBHCs on the power density spectra and the phase and time lags between different energy bands were discussed. From the comparison of the measured Fourier-frequency dependent time lags between two energy channels of the signal from Cyg~X-1, they deduced that a radial density dependence $n(r) \\propto r^{-1}$ in the Comptonizing hot plasma is appropriate to account for the observed hard time lags and the hard X-ray spectrum at the same time. However, their work was restricted to a geometry with a central soft photon source surrounded by a spherical corona of hot Comptonizing plasma. Furthermore, the density was the only parameter allowed to vary radially. In this paper, we extend the investigation of the effects of Comptonization the rapid variability in X-ray binaries, discussing two fundamentally different source geometries and radial gradients of other physical parameters, such as the electron temperature. This is done primarily with Monte-Carlo simulations of instantaneous flares of soft radiation (described as a $\\delta$-function in time), being Comptonized by the X-ray emitting region. The generalization to the case of multiple, randomly distributed flares (shot-noise) is straightforward. The light curves resulting from the Monte-Carlo simulations are then Fourier transformed, and the PSD as well as the Fourier-frequency dependent phase and time lags between different energy bands are calculated. In Section 2, we describe our Monte-Carlo simulations leading to the energy-dependent light curves for the different geometrical situations, together with analytical approximations for these light curves. Their Fourier transforms and the corresponding time lags between different energy bands will be discussed in Section 3. In Section 4, we investigate specific differences between the PSDs and hard time lags resulting from different geometries and compare the predictions of both scenarios to the power spectra and hard time lags observed for some GBHCs. We will suggest further observational tests which could either confirm a preferred geometry, or rule out to the Comptonization scenario in general. ", + "conclusions": "We presented a systematic study on the predicted Fourier power spectra and hard time lags resulting from two fundamentally different geometries in Comptonization models for the hard X-ray emission of Galactic black-hole candidates: a flaring soft photon source located in the center of the Comptonizing region (e. g., a hot corona) and a soft photon source located outside a hot inner disk region. The different scenarios yield different predictions for the Fourier power spectra and fundamental differences in the photon-energy dependence of the Fourier PSD slope and of the slope describing the hard-time-lag versus Fourier-period relation. We found that Comptonization of flares of external soft radiation generally leads to a weak photon-energy dependence of the power spectrum and the hard time lag, with the power spectrum being marginally consistent with the PSD measured for GX~339-1, but significantly steeper than for Cyg~X-1. Comptonization of centrally injected soft photon flares, in turn, leads to a significant photon-energy dependence in the PSD power-law slope at high Fourier frequencies and photon-energy dependent deviations from linearity of the hard time lags as a function of Fourier period. With a density gradient $n(r) \\propto r^{-1}$, the simulated PSD for medium-energy X-rays is consistent with the slope observed in Cyg~X-1. The turnover in the Fourier-frequency dependent hard time lag curve poses a severe challenge to Comptonization models for the rapid aperiodic variability in GBHCs in general due to the resulting large size estimates. Detailed measurements of the energy-dependent properties of the rapid aperiodic GBHC X-ray variability are strongly encouraged. Using the diagnostic tools developed in this paper, they can shed light on the question of the geometry of black-hole accretion disks and the location of the Comptonizing region and may serve to rule out a specific geometry or even the Comptonization models as a mechanism involved in the rapid aperiodic X-ray variability of Galactic black-hole candidates in general." + }, + "9802/astro-ph9802006_arXiv.txt": { + "abstract": "s{ The case for a flat Cold Dark Matter model with a positive cosmological constant $\\Lambda$ has been recently strongly advocated by some theoreticians. In this paper we give the observers point of view to the light of the most recent observations with a special emphasis on lensing tests. We confirm the apparent cosmic concordance for a flat Universe with $\\Omega_{\\Lambda}$ close to 0.6 but we note that a low mass density open universe with no cosmological constant is still quite acceptable for most of the reliable observational tests, including lensing tests as well.} ", + "introduction": "\\label{sec:intro} A model of the Universe should be consistent with the low mass density, as required by the cosmic evolution of the largest structures and nucleosynthesis, the observed value of the Hubble constant $H_0$, and an age compatible with the observations of globular clusters and with the ages of stellar populations in distant galaxies \\cite{Dunlopetal1996}. With these constraints the most likely region that survives in the space of cosmological parameters ($H_0$, $\\Omega_0$, $\\Omega_{\\Lambda}$) seems to force the evidence that we may live in a flat Cold Dark Matter Universe with a positive cosmological constant ($\\Lambda$-CDM). It is important to note that the recent measurements of the angular power spectrum of Cosmic Background Radiation (CBR) anisotropy are compatible with $H_0 \\simeq 65 \\ km/s/Mpc$ and $\\Omega_{\\Lambda} = \\Lambda/3 \\ H_0^2 \\simeq 0.6$ \\cite{Lineweaveretal1997}$^{\\!,\\,}$\\cite{Tegmark1996}. Indeed, projects like MAP and Planck Surveyor (2005-2007) that will allow a determination of the values of the three cosmological parameters with an accuracy of about 10\\% \\cite{Jungmanetal1996} or better \\cite{Bondetal1997} should in principle definitely close one of the most crucial debate in modern astrophysical cosmology. To follow the striking wording of Ostriker \\& Steinhardt \\cite{OstrikerandSteinhardt1996}, we have at present a \"Cosmic Concordance\" of data and models supporting a flat Universe with $\\Omega_{\\Lambda} = 0.63 \\pm 0.1$ . In fact, this Cosmic Concordance was first noted about two decades ago by Gunn \\& Tinsley \\cite{GunnandTinsley1975} and since then further re-investigated from time to time to the light of new observational data. More recently, Turner \\cite{Turner1997} and Bagla et al. \\cite{Baglaetal1996} presented up-dated discussions of the problem. They emphasize that there is still now no definitive theoretical arguments against a non-zero cosmological constant and that it is quite possible that we have to face all the consequences of such a result. The analysis of the domain of existence of a cosmological constant is possible because for Friedmann-Lemaitre 's models of the Universe the 3 fundamental cosmological parameters are linked to the curvature parameter $\\Omega_k = 1-\\Omega_0-\\Omega_{\\Lambda}$, where $\\Omega_0={8 \\pi G \\rho_0 \\over 3 H_0^2}$ is the density parameter expressed with the present-day density of the Universe, $\\rho_0$, and $\\Omega_{\\Lambda}$ is the reduced cosmological constant which appears as an additional term of energy. Excluding for the moment the results of gravitational lensing tests, let us summarize briefly the present state knowledge on $H_0$ and $\\Omega_0$, to the light of the latest observations. For the first time, all determinations of the value of the Hubble constant deduced from the HST observations of Cepheids are converging to a value around $63 \\pm 10 \\ km/s/Mpc$ (Sandage \\& Tammann \\cite{SandageandTammann1996} proposed $62 \\pm 6 \\ km/s/Mpc$). Beaulieu et al. \\cite{Beaulieuetal1997} have shown with a large sample of magellanic Cepheids that metallicity effects could increase this value up to $70 \\ km/s/Mpc$. Gouguenheim (this proceeding) presented a comprehensive discussion of $H_0$, including this problem and the Hipparcos contribution to the primary stellar distance scale. Moreover, despite an intrinsic limitation of lens modeling due to possible external gravitational shear or unseen additional masses on any line of sight, the time light delay in the two multiple quasars Q 0957+151 and PG 1115+080 also give the same result with about the same error bars \\cite{Kundicetal1996}$^{\\!,\\,}$\\cite{Falcoetal1997}$^{\\!,\\,}$\\cite{Keetoneta1996}. Last but not least, all others upper bound values coming from tests like absolute brightness of distant supernovae \\cite{Wheeleretal1997} and galaxies surface brightness \\cite{Thomsen1997} are fully compatible with this determination. If we consider now that $H_0$ is a least known with an accuracy of about 15\\%, almost all determinations of the density parameter $\\Omega_0$ still have large uncertainties ($0.2 < \\Omega_0 < 1$). By chance, the number of independent observational constraints on $\\Omega_0$ is large, so the peak of measurements around $0.3^{+0.2}_{-0.1}$ obtained from a large class of tests seems significant. Let us give a rapid overview of the situation for these major recent tests. Whatever the cosmic scenario of formation of large structures and galaxies, very high resolution N-body simulations always show that dark halos can be fitted with the same radial simple profile only scaled by two parameters. Navarro\\cite{Navarro1996} shows that the observed rotation curves of galaxies immersed in such Universal Dark Halos can only be explained for less concentrated halos such as those which are formed with a flat universe with $\\Omega_{\\Lambda}= 0.7$. The study of the redshift distribution of damped Lyman-$\\alpha$ halo absorbers has made important progresses during the last years. It shows that $\\Omega_{\\Lambda}= 0$ open-universes with $\\Omega_0< 0.4$ are inconsistent with the observations while a flat-universe with $\\Omega_{\\Lambda}= 0.6$ or with $\\Omega_0= 1$ \\cite{Gardner1997} are acceptable. Many others good constraints of the value of $\\Omega_0$ come from the study of masses content in clusters of galaxies, the largest (well) observed structures in the Universe. Their baryon fraction inferred from X-ray studies allows to deduce that $\\Omega_0 = 0.32 \\pm 0.2$ \\cite{Whiteetal1993} and their galaxies velocity fields give imply that $\\Omega_0 = 0.3 \\pm 0.1$ \\cite{Carlbergetal1997}, whereas the direct lensing measurements of their total gravitational masses within a radius of 0.5 Mpc gives $\\Omega_0$ in the range = 0.2-0.5 (cf Mellier et al, this conference). With the weak lensing method, higher lensing densities corresponding to $\\Omega_0$ close to 1 cannot be excluded from some tentative measurements of extremely weak gravitational shear in volume of several Mpc-sizes around cluster centers (Mellier et al., this proceeding). But for such observations the correction of instrumental distortion may not yet be fully under control \\cite{bm95}$^{\\!,\\,}$\\cite{kaiseretal95}. Surprisingly the lensing results obtained for large scale structures seems comparable to those coming from the latest interpretation of large cosmic flows \\cite{Deckel1997}. Note that the lensing method seems more reliable and is able to progress significantly in the near future. Therefore, the most likely value for $\\Omega_0$ emerging from all these observational constraints seems equal to $0.3^{+0.2}_{-0.1}$. At least, one crucial tests on $\\Omega_{\\Lambda}$ for flat models is now emerging from the observations of high-z supernovae. It allows in principle the determination of the acceleration parameter $q_o= \\Omega/2-\\Omega_{\\Lambda}/2$ with an accuracy of $\\pm$0.2 as soon as we will be able to detect and observe a sample of SNs at $z > 1$ \\cite{GoodbarandPerlmutter1995}. So far, with a sample of SNs only at $z < 0.4$ Perlmutter et al. \\cite{Pelmutteretal1996} note that $q_0$ could be actually in the 0-0.5 interval. Similarly, Bender et al. \\cite{Benderetal} recently demonstrate that we can actually use cluster ellipticals as cosmological candles. With two clusters at $z=0.375$ they tentatively found $q_0$ in the range 0-0.7. It is clear that such new observational tests do not much favor the Cosmic Concordance paradigm. They can progress a lot if suitable telescope time is given to these programs and we are expecting a lot from them for a direct measurement of $q_0$. ", + "conclusions": "If we accept that $H_0$ is around $60 \\ km/s/Mpc$ and that the Universe is old as it is inferred from the theory of the evolution of stars and stellar populations in very distant galaxies, it is impossible to have a flat matter-dominated Universe $\\Omega_0 =1$. In fact, many other observational tests including lensing tests also exclude this possibility except for a few that are not yet fully reliable because of possible observational bias or controversial theoretical hypothesis. It is true that many tests on the cosmological constant yield apparently concordant results that are in favor of a flat Universe with $\\Omega_0 = 0.3^{+0.2}_{-0.1}$ and $\\Omega_{\\Lambda} = 0.6 \\pm 0.2$. This is compatible with the reliable upper limit $\\Omega_{\\Lambda} < 0.65$ derived by Kochanek from quasars statistics. However, for most of the tests, open universes having $\\Omega_0 \\simeq 0.3$ and $\\Omega_{\\Lambda} = 0$ are also often acceptable. If we are actually living in a universe with a cosmological constant close to the Kochanek limit it is worth to improve the lensing tests particularly with the gravitational arcs. It is sure that the question will be also quickly addressed with the on-going large distant supernova surveys because the $q_0$ test has a remarkable efficiency. There is no doubt that the search for a non-zero cosmological constant corresponds to a high priority observing program without waiting for the MAP and Planck Surveyor surveys. At present and from the observational point of view it is still too early to make a case for a non-zero cosmological constant. More reliable observational facts have to be accumulated." + }, + "9802/astro-ph9802230_arXiv.txt": { + "abstract": "Observations of the high redshift Universe, interpreted in the context of a new generation of computer simulated model Universes, are providing new insights into the processes by which galaxies and quasars form and evolve, as well as the relationship between the formation of virialized, star-forming systems and the evolution of the intergalactic medium. We describe our recent measurements of the star-formation rates, stellar populations, and structure of galaxies and protogalactic fragments at $z\\sim$2.5, including narrow-band imaging in the near-IR, IR spectroscopy, and deep imaging from the ground and from space, using $HST$ and $ISO$. ", + "introduction": "Observations of the Hubble Deep Field, and other surveys of high redshift galaxies described in these proceedings, are contributing to a new picture of how large galaxies such as the Milky Way were assembled. One interpretation of the data so far is that large galaxies collapse out of what appears as several star-forming proto-galactic fragments at $z\\sim 2-3$ (Pascarelle et al. 1996; Haehnelt, Steinmetz $\\&$ Rauch 1998). Here we describe our search for H$\\alpha$ emission from the star-forming regions in high redshift galaxies, which is redshifted into the near-infrared. We focus on one of the outstanding questions about these objects: is the burst of star-formation seen in the UV continuum producing the dominant stellar population (by mass)? We show that for most high redshift objects, the data available to date provide only weak constraints on the age of the dominant stellar population (see also the discussion by Sawicki $\\&$ Yee 1998). An exception is our observation of the observed mid-IR (rest-frame near-IR) continuum of the $z=$2.7 galaxy MS1512-cB58 using $ISOCAM$. For this object, we can put a strong limit on the fraction of the mass comprised of an old population of stars. ", + "conclusions": "" + }, + "9802/astro-ph9802189_arXiv.txt": { + "abstract": "We investigate the metal contents of \\lya clouds at $2.210^{14.5}$ \\cm2 for which a mean metallicity of [C/H]$\\simeq -2.5$ was deduced. It has been suggested that the metals seen in these clouds may have been produced by a generation of Population III stars occurred at a much earlier epoch before the formation of quasars and normal galaxies. We shift the quasar spectra into the rest frame of each \\lya cloud with \\nrange and then average the rest-frame spectra in the spectral region encompassing the \\civ $\\lambda\\lambda$1548.20, 1550.77 absorption lines in order to vastly improve the S/N of the final composite spectrum. After eliminating a small number of \\lya lines that are estimated to have \\nrange which individually show detectable \\civ absorption, the remaining $\\sim 300$ lines whose corresponding \\civ $\\lambda$1548.20 regions are clean are divided into various samples to form composite spectra. No significant \\civ absorption is detected in any of the composite spectra we investigated. We derive an upper limit of $<0.088$ m\\AA\\ (95.5\\% confidence limit) for the equivalent width of the \\civ $\\lambda$1548.20 line, corresponding to [C/H]$<-3.5$ using the cosmological simulations of Rauch, Haehnelt, \\& Steinmetz (1997) to infer the ionization correction; the same simulation results have been used to deduce a mean metallicity of [C/H]$\\simeq -2.5$ for \\lya clouds with \\nhi$>10^{14.5}$ \\cm2. The mean metallicities of \\lya clouds with \\nrange, [C/H]$<-3.5$, is a factor of 10 lower than that inferred for \\lya clouds with \\nhi$>10^{14.5}$ \\cm2. This suggests that a sharp drop in the metallicity level of the intergalactic gas sets in at $10^{14}<$\\nhi$<10^{14.5}$ \\cm2. The result rules out the suggestion that a generation of Pop III stars could have polluted the entire universe to a (nearly) uniform metallicity level of [C/H]$\\simeq -2.5$. Cosmological simulations involving gas hydrodynamics indicate that \\lya clouds with \\nhi$>10^{14.5}$ \\cm2 mostly occur in the filamentary gaseous regions surrounding and connecting collapsed objects, while those with \\nhi$<10^{14}$ \\cm2 are preferentially found in void regions further away from collapsed structures. These results, coupled with the simulation results of Ostriker \\& Gnedin (1996) and Gnedin \\& Ostriker (1997) for Pop III star formation and enrichment, strongly suggest that most of the heavy elements in \\lya clouds with \\nhi$>10^{14.5}$ \\cm2 (i.e., gas in the filaments) were probably produced {\\it in situ} by Pop II stars, in the sense that they were either made by stars within the clouds or were ejected from nearby star-forming galaxies. Within this context, clouds with \\nhi$<10^{14}$ \\cm2 (i.e., gas in the void regions) may only have experienced pollution from Pop III stars. We derive constraints on the metallicities of the small number of \\lya clouds with \\nrange that individually show detectable \\civ absorption and discuss their possible nature. ", + "introduction": "The \\lya forest clouds seen in spectra of quasars, which are crudely defined to have \\nhi$<10^{16}-10^{17}$ \\cm2, were thought initially to contain primordial gas since no metal absorption lines were detected at their redshifts (Sargent et al 1980). However, moderate-resolution spectra taken with much higher S/N ratios (Meyer \\& York 1987) provided the first evidence that some of these clouds, especially those with relative large equivalent width, do show weak \\civ absorption. By shifting to the rest frame and averaging the \\civ regions of hundreds of \\lya lines which did not show \\civ absorption individually, Lu (1991) provided a probably detection of \\civ absorption associated with these clouds at the $w_r(1548)\\sim 7$ m\\AA\\ level. Subsequent high resolution, high S/N spectra taken with the 10m Keck telescopes (Cowie et al 1995; Tytler et al 1995; Songaila \\& Cowie 1996) have clearly revealed (mostly) weak \\civ absorption associated with individual \\lya clouds at $z\\sim 3$ with \\nhi\\ as low as $10^{14.5}$ \\cm2. Specifically, it was found that essentially all clouds with \\nhi$\\geq 10^{15}$ \\cm2 show \\civ absorption, and that 50-75\\% of the \\lya clouds with \\nhi$\\geq 10^{14.5}$ \\cm2 show \\civ absorption. The presence of \\civ absorption in \\lya clouds with \\nhi$> 10^{14.5}$ \\cm2 indicates that these clouds have been enriched in heavy elements at redshifts as high as $z>3.6$. Recent cosmological simulations of structure formation involving gas processes (hydrodynamics) suggest that the \\lya lines at $z>2$ represent low density gas in-between collapsed objects and are closely related to the formation of galaxies, clusters, voids, and other structures in the universe (Cen et al 1994; Petitjean, Mucket, \\& Kates 1995; Zhang, Anninos, \\& Norman 1995; Hernquist et al 1996; Miralda-Escude et al 1996). An important issue is whether the metals seen in these intergalactic clouds were produced locally (i.e., produced by stars formed within the clouds, or ejected from nearby star-forming galaxies) or were produced at a much earlier epoch in a wide-spread fashion (i.e., by Population III stars; Ostriker \\& Gendin 1996). If the metals in \\lya clouds were produced {\\it in situ}, one may expect that clouds in regions of space sufficiently far away from star-forming objects should be metal-free. Whereas in the case of Pop III star enrichment, most of the universe may have been polluted to some degree and few clouds may be entirely metal-free. While it is not yet possible to directly see the spatial distributions of star-forming objects and \\lya clouds at high redshift, some insight is provided by cosmological simulations. These simulations indicate that clouds with $10^{14.5}<$\\nhi$<10^{17}$ \\cm2 account for 40-60 \\% of the total baryons in the universe at $z=2-4$ (Miralda-Escude et al 1996), and that such clouds preferentially occupy the filamentary gaseous regions surrounding and connecting collapsed structures (i.e., galaxies). The simulations also indicate that clouds with \\nhi$<10^{14}$ \\cm2 predominantly occupy the void regions. These results lead to the idea that one can perhaps obtain clues about the origin of the metals by investigating the metal contents of the very low column density (i.e., \\nhi$<10^{14}$ \\cm2) clouds: the presence of any significant amount of metals in these clouds would favor the Pop III idea for the reasons discussed above. In this paper we present the results of a search for \\civ absorption associated with \\lya clouds with $10^{13.5}<$\\nhi$<10^{14}$ \\cm2, using high resolution, high S/N quasar spectra obtained with the 10m Keck I telescope and the High Resolution Spectrometer (HIRES; Vogt 1992). Given the low \\nhi\\ of the clouds, it would be practically impossible to obtain spectra of quasars with sufficient S/N to detect the \\civ absorption in individual clouds even if the N(\\civ)/\\nhi\\ ratio remains the same as that in clouds with higher \\nhi\\ (S/N at least a factor of 5 better than the current best spectrum would be required). The problem is compounded by the fact that the N(\\civ)/\\nhi\\ ratio is predicted to drop with decreasing \\nhi\\ at \\nhi$<10^{14}$ \\cm2 as the clouds become more ionized (e.g. Rauch, Haehnelt, \\& Steinmetz 1997). Accordingly, the technique we adopt is the composite spectrum method pioneered by Norris, Hartwick, \\& Peterson (1983). Essentially, one shifts a quasar spectrum into the rest frame of each \\lya absorption line and then averages the rest-frame spectra over the spectral region encompassing the \\civ $\\lambda\\lambda$1548.20, 1550.77 lines to form a composite spectrum. The S/N in the composite spectrum increases roughly as the square root of the number of lines being averaged, making this a very effective way to achieve ultra-high S/N without using unrealistically large amounts of telescope time. The end result is then an average \\civ absorption associated with an ensemble of \\lya clouds. This technique has been applied previously to \\lya clouds to search for \\civ absorption by Williger et al (1989), Lu (1991), Tytler \\& Fan (1994), and Barlow \\& Tytler (1998), and to metal systems to search for O VI absorption by Lu \\& Savage (1993). Incidentally, this averaging process automatically suppresses any residual noise features in the spectra caused by imperfect flatfielding or other artifacts and is ideal for achieving ultra-high S/N in the averaged spectra. In section 2 we describe the spectra used in the analysis and the procedure for selecting \\lya lines with the desired column densities (i.e., \\nrange). The composite spectra for various samples of \\lya clouds are discussed in section 3. No \\civ absorption is detected in any of the composite spectra, thus yielding only upper limits to the metallicity of the clouds. In section 4 we discuss the implications of our results for the nature of the metal enrichment in the \\lya clouds. We briefly summarize our main conclusions in section 5. ", + "conclusions": "We have investigated the metal contents of \\lya clouds at $2.210^{14.5}$ \\cm2 for which a mean metallicity of [C/H]$\\simeq -2.5$ was deduced. It has been suggested that the metals seen in these clouds may have been produced by a generation of Population III stars occurred at a much earlier epoch before the formation of quasars and normal galaxies. From 9 of our best S/N quasar spectra obtained with the HIRES, we selected $\\sim$300 \\lya lines with estimated \\nhi\\ between \\nrange. The \\lya lines selected have absorption redshift $2.210^{14.5}$ \\cm2. The low metallicities derived for the \\nrange clouds appear to hold to within a factor of $\\sim 2$ independent of redshift or \\nhi\\ over the range considered above. In particular, a sample of \\lya clouds occurring within 200 \\kms of known metal systems yielded a mean metallicity of [C/H]$<-3.0$ (95.5\\% confidence limit) based on a composite spectrum analysis. (2) The mean metallicity of \\lya clouds with \\nrange, [C/H]$<-3.5$, is a factor of 10 lower than that inferred for \\lya clouds with \\nhi$>10^{14.5}$ \\cm2. This suggests that a sharp drop in the metallicity level of \\lya\\ clouds sets in at \\nhi$=10^{14}-10^{14.5}$ \\cm2. This result rules out the suggestion that a generation of Pop III stars could have polluted the entire universe to a uniform level of [C/H]$\\simeq -2.5$. Cosmological simulations involving gas hydrodynamics indicate that \\lya absorption with \\nhi$>10^{14.5}$ \\cm2 mostly occurs in continuous filaments of gas surrounding and connecting collapsed objects, while those with \\nhi$<10^{14}$ \\cm2 are preferentially found in void regions further away from collapsed objects. These results, coupled with theoretical predictions about Pop III star formation and enrichment (Ostriker \\& Gnedin 1996; Gnedin \\& Ostriker 1997), strongly suggest that most of the heavy elements in \\lya clouds with \\nhi$>10^{14.5}$ \\cm2 were probably produced {\\it in situ} by Pop II stars, in the sense that they were either made by stars within the clouds or were ejected from nearby star-forming galaxies. Consequently, the void regions (\\nhi$<10^{14}$ \\cm2) could only have experienced the enrichment by Pop III stars. We estimated a typical thickness of the order of 100 kpc at $z=3$ for the metal-enriched filamentary structures. (3) The low metallicities we inferred for the \\lya clouds with \\nrange require that, if there was a generation of Pop III stars that occurred in the very early universe, the mean enrichment level resulting from the explosions of such stars should be [C/H]$<-3.5$ for \\lya clouds with \\nrange at $2.210,000$ \\kms away from the emission redshift of the background quasars and are likely to be intervening in nature. Collisional ionization is ruled out for at least 3 of the 7 intervening clouds based on the width of the \\civ absorption lines. If the clouds are photoionized by the integrated light from quasars, as is commonly assumed, the lower limits to their [C/H] values are inferred to be between $-2.2$ and $-0.9$. The relatively high metallicities of these clouds are in stark contrast with those of other similar \\nhi\\ clouds that do not show \\civ absorption individually for which we inferred a mean [C/H]$<-3.5$ from their composite spectrum. If \\lya clouds with \\nrange have a metallicity distribution that can be approximated by a Gaussian, and if these clouds that individually show detectable \\civ absorption represent the high metallicity tail of this distribution, we infer that the metallicity distribution should have a mean [C/H]$<-4$ with a dispersion of at least 1 dex. Alternatively, the \\lya clouds which individually show detectable \\civ absorption may have a separate origin." + }, + "9802/astro-ph9802140_arXiv.txt": { + "abstract": "We present mass distributions obtained from near-infrared (NIR) surface brightness decompositions and rotation curve fitting of a sample of early-type spiral galaxies. Bulge and disk mass-to-light (M/L) ratios, dark halo parameters, and the modified Newtonian dynamics (MOND) acceleration parameter are derived. We find that the mean disk NIR M/L is higher than that of the bulge, and comparison with stellar population synthesis models implies that early-type spiral bulges are, on average, younger and more metal rich than disks. NIR disk M/L is found to depend on disk luminosity, consistently with previously reported trends in ellipticals and spirals, and with cold dark matter models for disk formation. Within the optical radius, the mean ratio of stellar to dark matter is 2 and the typical dark halo mass is 10$^{11}$\\,M$_\\odot$. The value of the MOND acceleration parameter that best accommodates the sample as a whole is $1.3\\cdot10^{-8}$~cm\\,s$^{-2}$. ", + "introduction": "The comparison of the luminosity profiles and rotation curves of disk galaxies provided the first evidence for dark matter in outer parts of galaxies (e.g., van Albada \\& Sancisi \\cite{vas}). The luminous material, either in the form of stars or neutral gas, is not able to reproduce the approximate flatness of the rotation curves at large radii, implying that the Newtonian theory of gravitation needs to be modified or that the mass in the outer regions of most spirals is dominated by a dark component. The physical extent of such a ``dark halo'' cannot be inferred from the rotation curves, since these last are flat even where the luminous material ceases to be detected. Hence, the total mass in spiral galaxies is not yet well known. The properties of dark matter (DM) and the structure of dark halos as derived from rotation curves have profound implications on cosmological issues such as galaxy formation. Recent N-body simulations of cold dark matter halos together with adiabatic infall models have shown that matching observed rotation curves requires a systematic increase of disk mass-to-light ratios with luminosity (Navarro et al. \\cite{navarro}). Such a trend is observed in large bodies of rotation curve data (Broeils \\cite{broeils}; Persic et al. \\cite{pss} -- hereafter PSS), and further comparisons between theoretical predictions and observations may help clarify the nature of dark matter in galaxies. The presence of a dark halo makes global mass-to-light (M/L) ratios derived from rotation curves much higher than those expected from stellar populations, and produces radial gradients in the total M/L. The contribution of DM tends to be larger for low-luminosity systems, as the fraction of DM to luminous matter within the optical radius appears to decrease with increasing galaxy luminosity (Kormendy \\cite{kormendy90}; Salucci et al. \\cite{sap}). In luminous systems, therefore, the mass distribution in the inner few disk scale lengths can be largely ascribed to the bulge and disk stars. Consequently, when stellar galaxy components can be accurately isolated with surface photometry, estimates of their M/L from rotation curves provide useful constraints for the age, abundance, and star formation history of their stellar populations. Recent population synthesis models predict trends of M/L with metallicity and with wavelength (Worthey 1994; Bruzual \\& Charlot 1993), and can be used to infer and compare properties of the stellar populations in bulges and disks. We have reported in a previous paper (Moriondo et al. \\cite{moriondo} -- hereafter Paper I) the results of two-dimensional near-infrared (NIR) surface brightness decompositions for a sample of early-type spirals. NIR wavelengths, especially when combined with those in the optical, are a powerful diagnostic tool since they trace more accurately the stellar mass content and minimize the complications of extinction. In this paper, we apply the bulge$+$disk decomposition results to the analysis of the rotation curves for our sample. In Sect.~2 we derive the radial mass distributions and evaluate the contribution of the bulge, disk, and dark halo to the observed rotation curves; as an alternative to dark matter halos, we also derive a value of the modified Newtonian dynamics critical acceleration parameter. The resulting NIR M/L ratios of the luminous components are analyzed in Sect.~3, and compared with those at optical wavelengths and with population synthesis models. Finally, we assess the importance of dark halos, and explore trends of M/L with luminosity in the context of the properties of the fundamental plane for gravitationally bound objects (e.g., Burstein et al. \\cite{burstein97}). ", + "conclusions": "\\subsection{Stellar content}\\label{section:sps}\\protect We have compared colors and M/L's for our sample with models of stellar population synthesis (SPS) to place constraints on age and metallicity of the average stellar content of bulges and disks. In particular we have considered Worthey's model (\\cite{worthey}, W94 hereafter) for single stellar populations with ages from 1.5 to 17 Gyr and different metallicities, and the 1995 release of the model by Bruzual and Charlot (\\cite{br_ch}, BC95) for populations with solar metallicity and different star formation histories. The mean colors and M/L's of bulges and disks are plotted in Fig. \\ref{figure:models} together with the two SPS models. The average M/L for the bulges are computed excluding values below 0.2 in all bands. W94 colors and M/L's agree rather well with the observed values, whereas BC95 have both bluer colors and lower M/L's. The discrepancies between the two models have been discussed by Charlot et al. (\\cite{cwb}). Inspection of the left panel of Fig. \\ref{figure:models} shows that bulges are notably redder than disks, both in $r-K$ and in $J-K$ (see also Paper I). Nevertheless, the discrepancies between different SPS models, the degeneracy age/metallicity, and the possible effect of extinction make the color-color plot in Fig. \\ref{figure:models} ambiguous as a diagnostic for distinguishing different stellar populations. \\begin{figure*} \\picplace{1.5cm} \\caption[]{Left panel: Comparison of average component colors to models of stellar population synthesis. The three thick lines with marked dots correspond to models (W94) with different metallicities as labelled: [Fe/H]=0.5 [Fe/H]=0.0, and [Fe/H]=-0.25. Age increases from bottom to top; designated points correspond to ages of 1.5, 2, 3, 5, 8, 12, 17 Gyr. The three thin continuous lines correspond to models (BC95) with fixed solar metallicity and different IMF's. Dots mark the same ages as for W94. models. Filled symbols represent the average colors of disks with the relative uncertainties; the triangle is for the parametric results and the square for the non parametric ones. Empty symbols are the bulge values. Right panel: the same comparison for M/L's ($r$ and $K$ bands).} \\label{figure:models}\\protect \\end{figure*} Fortunately, M/L ratios disentangle, at least partly, the ambiguity of age and metallicity. We find that bulges, on average, have lower M/L's than disks in all bands\\footnote{$J$ is not shown in the Fig. \\ref{figure:models}, but $J$-band M/L's behave in the same way as those in $r$ and $K$.}. Moreover, the values of M/L seen in the right panel of Fig. \\ref{figure:models}, according to the SPS predictions, suggest that bulges are {\\it younger} and {\\it more metal rich} than disks. We note that both W94 and BC95 models of a given abundance follow approximately the same trend with age. Hence, the displacement of the bulge and disk values relative to that trend implies that bulges are characterized by a younger age than that of the disks, independently of discrepancies between models. Except for extreme inclinations, $i > 75^{\\circ}$, dust in the central regions affects the bulge more than the disk (Bianchi et al. \\cite{bianchi}). Consequently, the noted difference between bulge and disk M/L cannot be attributed to internal extinction, since the correction would only increase the difference between the two. The possibility that the disk M/L's have been overestimated seems rather unlikely, even under the MBD hypothesis (although see Bottema \\cite{bottema}; Courteau \\& Rix \\cite{courteau:rix}) since an average error greater than 30\\% on the estimated disk RC contribution would be required to make the disks' average age comparable to bulges'. Finally, the bulge M/L's could have been systematically underestimated. This would be the case only if all the RC's (and not just the few already noted and excluded from the mean) rose too slowly in the inner regions with respect to the true circular velocity of the galaxy. We discussed briefly this point in Sect.~\\ref{results:m_l}, concluding that most likely such underestimates, if present, are not sufficient to eliminate the observed difference in M/L between the components. That bulges may be more metal rich than disks is not a new result (Bica \\& Alloin \\cite{ba}; Delisle \\& Hardy \\cite{dh}; Giovanardi \\& Hunt \\cite{giova}; Paper I), and abundance variations are thought to be driven by variations with mass (e.g., Zaritsky et al. \\cite{zaritsky}). That bulges appear to be {\\it younger} than disks is somewhat surprising; the comparison with SPS models shown in Fig. \\ref{figure:models} implies an age difference of around 50\\%, or 5~Gyr. Nevertheless, such a result may be interpreted in light of recent observational and theoretical work on bulge dynamics. Many bulges show kinematic and photometric signatures usually associated with disks, including flattened distributions, exponential fall-off, dominance of the rotation velocity component, and spiral structure in the bulge-dominated region (Kormendy \\cite{kormendy93} and references therein). Moreover, some bulges have blue colors, the result of extremely young populations (Schweizer \\cite{schweizer}), and as noted in Paper I, at least three of the galaxies in our sample appear to be actively forming stars\\footnote{Two of these have, however, been excluded from the calculation of the mean.}. As suggested by Kormendy and others, ``bulges'' may be built up over time from disk material transported to the central regions by gravitational perturbations; such bulges would appear younger than the disks from which they derive. \\subsection{Correlations with mass-to-light ratios}\\protect\\label{ml} Since, as for the photometric properties discussed in Paper~I, disk characteristics are more reliably determined than those of bulges, we will concentrate on the M/L's obtained for the disks. Several authors (see for instance Djorgovski \\& Santiago \\cite{ds}, and references therein) have demonstrated that elliptical galaxies follow a relation which can be expressed in terms of a power law: M/L~$\\propto L^{\\alpha}$ with $\\alpha=0.2\\sim0.3$ depending on the sample and the photometric band. Kent (\\cite{kent86}) also found that the disks of spirals follow a similar relation in the $r$ band at fixed morphological type with an exponent $\\alpha\\,=\\,0.18\\,\\pm\\,0.07$. More recently Persic, Salucci, and collaborators (PSS; Salucci et al. \\cite{sap}) found for a sample of late-type spirals that in $B$, $\\alpha = 0.30 \\sim 0.35$. Burstein et al. (\\cite{burstein97}) have suggested that a relation of this kind between M/L and luminosity is common to all the self-gravitating structures in the universe, ranging from globular clusters to clusters of galaxies. Finally, models of cold dark matter halos, based on N-body simulations and adiabatic infall for disk formation, require a variation of disk M/L with $B$ luminosity in order to accommodate observed rotation curves (Navarro et al. \\cite{navarro}). \\begin{figure*} \\picplace{1.5cm} \\caption[]{The M/L vs luminosity for the disks ($K$ band) with the trend derived from $\\alpha_B$ and the color-luminosity relation described in the text. } \\label{figure:ml_l}\\protect \\end{figure*} We have investigated the compatibility of the values of $\\alpha$ found in the optical with our $K$-band data, assuming the trend of M/L with $L$ is due to disk stars and not to DM. We can convert the index $\\alpha_B$ found in the $B$ band to the $K$-band value using the well-established color--luminosity relation for spirals (Visvanathan \\cite{vis}; Wyse \\cite{wyse}; Tully et al. \\cite{tma}). Based on recent data, Gavazzi (\\cite{gavazzi}) finds: \\[B-H = -0.2\\, {\\cal M}_H + const.\\] for early-type spirals. If we assume that $H-K$ is independent of luminosity (which is likely since $H-K$ is typically small, $\\sim$~0.2~mag), then we have a similar relation for $B-K$ with \\[B-K = -0.2\\, {\\cal M}_K + const. \\;\\; . \\] We take $\\alpha_B\\,=\\,$0.35 (PSS) together with the slope of the color--luminosity relation defined above, we infer a value of $\\alpha_K\\,=\\,0.15$. If we {\\it fit} our data to a regression of M/L versus disk luminosity, we obtain $\\alpha_K\\,=\\,0.22\\,\\pm\\,0.15$, consistent with the expected value of 0.15. Figure \\ref{figure:ml_l} shows $K$-band disk luminosity plotted against disk M/L ($K$), together with a line having slope $\\alpha\\,=\\,0.15$. The M/L vs luminosity relation for elliptical galaxies is implicit in a more general relation, namely the one defining the fundamental plane (FP) of elliptical galaxies (see review by Kormendy \\& Djorgovski \\cite{kd}): \\begin{equation} R_e \\propto \\sigma_c^a I_e^b \\label{eqn:fp}\\protect \\end{equation} where $R_e$ and $I_e$ are respectively the effective radius and surface brightness and $\\sigma_c$ is the observed central velocity dispersion\\footnote{We note that for elliptical galaxies the central observed velocity dispersion depends on {\\it total} mass, whereas RC fitting in principle disentangles the contributions of stars and DM. If we assume, following Salucci \\& Persic \\cite{sal_per}, that the fractional DM content is a decreasing function of luminosity also for ellipticals, an M/L vs $L$ relation for the stellar component alone would require a steeper slope than the one derived from the FP relation.}. We have fitted a similar relation between disk scale lengths, central brightnesses, and peak rotation velocities of our disks, and find: $a=1.2 \\pm 0.2$ and $b=-0.84 \\pm 0.1$. This implies that the disks of early-type spirals also define a plane similar to that for ellipticals; the elliptical FP has $a=1.4 \\pm 0.15$ and $b=-0.9 \\pm 0.1$. A similar result, but for the photometric properties only, was reported in Paper~I, both for bulges and disks. When we derive the M/L vs $L$ relation from Eq. \\ref{eqn:fp} and the virial theorem we obtain: \\[ \\mbox{M/L} \\propto L^{\\frac{2-a}{2a}}I(0)^{-\\frac{a+4b+2}{2a}} \\] which contains a residual dependence on the central brightness $I(0)$. However our values for $a$ and $b$ yield \\[ \\mbox{M/L} \\propto L^{0.3}I(0)^{-0.07} \\] where, as for ellipticals, the dominant dependence is on luminosity. As discussed by Djorgovski \\& Santiago (\\cite{ds}), a relation between M/L and luminosity (or mass) can come about in several ways. One possibility is that the disk M/L's are contaminated by a DM contribution which has a density profile similar to that of the stellar disk. The sense of the M/L vs $L$ relation would require this DM fraction to increase with luminosity, contrary to the trend observed for the global DM fraction which increases with {\\it decreasing} luminosity (e.g., PSS). Alternatively, the MBD hypothesis could be incorrect, and the stellar disk M/L constant. In this case, though, the trend in Fig. \\ref{figure:ml_l} would require the MBD hypothesis to be more valid in lower luminosity systems, contrary to common beliefs. Another possibility is that disks of different luminosities harbor different stellar populations, which is also suggested by the color-luminosity relation mentioned above. W94 predicts that at fixed age, initial mass function (IMF), and star formation rate (SFR), M/L is an increasing function of metallicity in the optical, but a {\\it decreasing} one in the NIR. This suggests that the M/L vs $L$ correlation cannot be understood in terms of a metallicity variation. Alternatively, such a correlation could be driven by a change of average age or star formation history with luminosity. Again, the observed difference in the slope of the correlation at different wavelengths can be compared to the predictions of SPS models. To test this possibility we made use of the BC95 models, at fixed (solar) metallicity and IMF (Salpeter \\cite{salpeter}), considering single burst populations at different age $T$, and populations with different e-folding time $\\tau$ (exponential SFR) at fixed age (10 Gyr). We have approximated the model dependence of M/L on $T$ and $\\tau$ with power laws, whose index in $B$ and $K$ has been determined with a best fit. Assuming M/L$_B \\propto L^{0.35}$, in the first case the model predicts a difference in the slope of this relation of 0.17 passing to the $K$ band, consistent with the predicted difference of $\\sim$~0.2. Considering the variation with $\\tau$ the prediction is $\\sim$ 0.2, in similarly good agreement. It, therefore, seems plausible that the variation of M/L with luminosity is driven by different star formation histories, and the consequent different stellar mixes. \\subsection{Dark halos}\\protect\\label{dh} We claim to recognize the presence of a dark halo in six of our galaxies: three of them are pseudo-isothermal and three constant-density spheres. We do not find any systematic difference between the two models, at least in terms of central densities or masses within $R_{25}$. The statistics are sparse, but we can attempt a comparison with the general relations found by PSS in the $B$ band for a large sample of late-type spirals. In particular, we can check that the ratios of the halo central density to the critical density, $\\rho_0/\\rho_c$, and the dark to visible mass at the optical radius, $M_h/M_V$, are consistent with the correlations with the $B$-band galaxy luminosity given in PSS. They find: \\[ \\frac{M_h}{M_V}=0.4\\left(\\frac{L_B}{L_B^*}\\right)^{-0.9} \\;\\;\\; \\mbox{and} \\;\\;\\;\\;\\; \\frac{\\rho_0}{\\rho_c}=3.5\\cdot 10^4\\left(\\frac{L_B}{L_B^*}\\right)^{-0.07} \\] with $L_B^*=10^{10.4}L^{\\odot}_{B}$. If we assume the same color--luminosity relation as in Sect.~\\ref{ml}, we find in the $K$ band after scaling to $H_0=50$ km\\,s$^{-1}$: \\[ \\frac{M_h}{M_V}=0.8\\left(\\frac{L_K}{L_K^*}\\right)^{-0.7} \\;\\;\\; \\mbox{and} \\;\\;\\;\\;\\; \\frac{\\rho_0}{\\rho_c}=3.7\\cdot 10^4\\left(\\frac{L_K}{L_K^*}\\right)^{-0.06} \\,\\, . \\] with $\\rho_c=3H_0^2/8\\pi G=0.5\\cdot 10^{-29}$ g\\, cm$^{-3}$ and ${\\cal M}^{*}_{K}=-23.97$, the absolute magnitude corresponding to $L_K^*$. We estimated ${\\cal M}^{*}_{K}$ from the color--luminosity relation, with the zero order coefficient fixed by the median values of ${\\cal M}_B$ and $B-K$ for our sample. The trends shown in Fig. \\ref{figure:halo} are consistent with the anticorrelation between dark and luminous mass found in PSS. Moreover, there is no striking discrepancy between our galaxies and the behavior of later type systems, suggesting that dark halos are similar for all spirals. In the right panel of Fig. \\ref{figure:halo} our data reveal roughly the same $M_h$ as found in late-type galaxies with the same $B$ luminosity. \\begin{figure*} \\picplace{1.5cm} \\caption[]{Central halo density and ratio of dark to visible mass vs galaxy $K$ absolute magnitude. Filled triangles correspond to isothermal halos, open circles to constant density ones. The dashed lines represent the NIR relations between these parameters corresponding to the ones obtained by PSS in the $B$ band.} \\label{figure:halo}\\protect \\end{figure*}" + }, + "9802/astro-ph9802154_arXiv.txt": { + "abstract": "The Gamma-Ray burst detector on {\\it Ginga} consisted of a proportional counter to observe the x-rays and a scintillation counter to observe the gamma-rays. It was ideally suited to study the x-rays associated with gamma-ray bursts (GRBs). {\\it Ginga} detected $\\sim120$ GRBs and 22 of them had sufficient statistics to determine spectra from 2 to 400 keV. Although the {\\it Ginga} and BATSE trigger criteria were very similar, the distribution of spectral parameters was different. {\\it Ginga} observed bend energies in the spectra down to 2 keV and had a larger fraction of bursts with low energy power law indexes greater than zero. The average ratio of energy in the x-ray band (2 to 10 keV) compared to the gamma-ray band (50 to 300 keV) was 24\\%. Some events had more energy in the x-ray band than in the gamma-ray band. One {\\it Ginga} event had a period of time preceding the gamma rays that was effectively pure x-ray emission. This x-ray ``preactivity'' might be due to the penchant for the GRB time structure to be broader at lower energy rather than a different physical process. The x-rays tend to rise {\\it and} fall slower than the gamma rays but they both tend to peak at about the same time. This argues against models involving the injection of relativistic electrons that cool by synchrotron radiation. ", + "introduction": "Gamma-ray bursts (GRBs) are aptly name, most of their power is usually emitted in the 100 keV to 1 MeV energy range. Early results indicated that only a few percent of the energy of a Gamma-Ray burst (GRB) occurs in the 2-to-10 keV x-ray range \\cite{laros84,katoh84}, although the X-ray emission might outlast the main GRB event in some bursts (so called X-ray tails). Based on these intriguing results, the gamma-ray burst detector (GBD) flown aboard the {\\it Ginga} satellite was specifically designed to investigate burst spectra in the X-ray regime \\cite{mur89}. {\\it Ginga} was launched in February of 1987, and the GBD was operational from March, 1987, until the reentry of the spacecraft in October, 1991. More recently, the BeppoSax satellite has observed several bursts in both x-rays and gamma-rays \\cite{piro97,piro98,frontera98} and has discovered soft x-ray afterglows \\cite{costa97}. This latter discovery has opened the way for the long-sought GRB counterparts, including one with a measured redshift \\cite{metz97}. Here, we summarize x-ray results from the {\\it Ginga} experiment. ", + "conclusions": "" + }, + "9802/astro-ph9802224_arXiv.txt": { + "abstract": "I review recent studies about the gravitational potential and stellar dynamics of M87 in particular, and the dynamics of the stars in the presence of a super-massive central black hole, in general. At large radii, investigations of both the X-ray emitting gas and the velocity distribution of globular clusters indicate the presence of large amounts of non-luminous matter, possibly belonging to the inner parts of the Virgo cluster. At small radii, there is no evidence from the stellar kinematics, at most a hint, for the existence of a central point mass, whereas the gas dynamics reveal the presence of a highly concentrated mass in the centre of M87, possibly a super-massive black hole (\\bh). Given the existence of such a central mass, the stellar kinematics indicate a strong tangential anisotropy of the stellar motion inside a few arcseconds. The implications of this result for the evolution and formation history of M87 and its central \\bh\\ are discussed. I also discuss in more general terms the structural changes that a highly concentrated central mass can induce in its parent galaxy. ", + "introduction": "\\index{galaxy formation} According to their observable properties, elliptical galaxies can be divided into two classes. This dichotomy is most clearly revealed in the central brightness distribution (cf.\\ Fig.~3 of Gebhardt \\etal\\ 1996). Consequently, the two classes are commonly called `power-law galaxies' and `core galaxies' (though -- as so often in astronomy -- the latter term is highly misleading: these galaxies actually do not have a core of constant density). Core galaxies have shallow central luminosity density profiles with $j\\propto r^{-\\gamma}$, $\\gamma\\la1.3$, ellipticities E0 to E3-4, elliptic to boxy isophotes, and negligible rotation $v_{\\rm rot}\\ll \\sigma$. They are on average bright ($M_V\\la-19.5$) and often radio-loud and X-ray-active, possess extended stellar envelopes and rich ($N\\ga2000$), extended globular cluster (GC) systems which are multi-modal in their properties. These galaxies are thought to be of round to triaxial shape supported by anisotropy in the stellar motions. Power-law galaxies, on the other hand, have steep central density cusps with $\\gamma\\ga-1.5$, ellipticities up to E7, elliptic to disky isophotes, and significant rotation velocities $v_{\\rm rot}\\sim\\sigma$. They are on average fainter ($M_V\\ga-21.5$) and show no radio or X-ray activity; their surface density follows a de Vaucouleurs profile and their GC systems are poor ($N\\la1500$) and with a profile following that of the stellar light. These galaxies are believed to be of near-oblate shape supported by rotational motions. Clearly, M87 having a shallow density cusp ($\\gamma\\approx1.2$), round isophotes, neglible rotation, radio and X-ray activity, an extended stellar envelope, and a rich and bi-modal globular cluster system is a generic representative of the class of core galaxies. It is generally believed, that the core galaxies are formed by one or more major merger events. In face of this hypothesis, it is important to ask whether M87 is consistent with being a merger remnant. \\begin{figure} \\centerline{ \\epsfxsize=8cm \\epsfbox{Dehnen.fig1.ps}} \\caption[]{Velocity dispersion for M87 as measured for stars (filled circles, van der Marel 1994) and planetary nebul\\ae\\ (open squares, Cohen \\& Ryzhov 1997). The full line is the rotational velocity measured by the latter authors.} \\label{fig-sigma} \\end{figure} \\begin{figure} \\centerline{ \\epsfxsize=8cm \\epsfbox{Dehnen.fig2.ps}} \\caption[]{Mass enclosed in radius $r$ as derived from studies of stellar kinematics (open circles, Sargent \\etal\\ 1978), globular clusters (triangle, Mould \\etal\\ 1990, and filled circles, Cohen \\& Ryzhov 1997), and X-ray gas (thin lines: lower and upper limits, Nulsen \\& B\\\"ohringer 1995, corrected for the difference in the adopted distance to M87). The solid line represents a power-law fit to the filled circles: ${\\rm M}\\propto r^{1.7}$.} \\label{fig-mass} \\end{figure} ", + "conclusions": "The kinematics of M87 are well studied, which make this galaxy a good test case for the theories of galaxy formation. Outside $\\sim100\\arcsec$, the velocity dispersion profile rises indicating the presence of large amounts of non-luminous matter. The inferred density profile $\\rho\\propto r^{-1.3}$ is consistent with predictions from CDM cosmogony for the inner parts of dark-matter halos. The massive black hole (\\bh), detected in the very centre of M87 by gas motions, together with the observed stellar kinematics implies a significant tangential anisotropy of the stellar motions. Among the formation histories discussed for a \\bh\\ in a galactic centre, only the model of accretion of other massive \\bh s, originating from the centres of cannabalized companions, can explain such a strong anisotropy. This scenario also predicts a shallow stellar density cusp as observed for M87. (Quantitatively, there are some discrepancies, which may well be due to over-simplification in the simulations of this process.) A massive \\bh\\ at the centre of a triaxial galaxy renders, by the destruction of box orbits, the shape of its host axisymmetric. This mechanism becomes very fast once the \\bh\\ mass reaches a critical value, which is of the order of 1\\% of its host's mass. Since the conservation of angular momentum along ballistic orbits in an axisymmetric galaxy obstructs gas-fueling of the centre, this process may pose an upper limit for the mass a \\bh\\ can reach by gas-accretion. An upper limit of this order is indeed observed among \\bh\\ masses inferred from the dynamics of early-type galaxies." + }, + "9802/astro-ph9802012_arXiv.txt": { + "abstract": "If the universe is multiply-connected and sufficiently small, then the last scattering surface wraps around the universe and intersects itself. Each circle of intersection appears as two distinct circles on the microwave sky. The present article shows how to use the matched circles to explicitly reconstruct the global topology of space. ", + "introduction": "If the universe is multiply-connected and sufficiently small, then the last scattering surface (LSS) wraps around the universe and intersects itself~\\cite{css0}. Each circle of intersection appears as two distinct circles on the microwave sky, even though the two images correspond to the same circle of points in space itself. In their article in this issue, Cornish, Spergel and Starkman~\\cite{css1} show how to find such pairs of matching circles from the high-resolution data to be provided by NASA's Microwave Anisotropy Probe (MAP) in the year 2001, or by the ESA's Planck satellite a few years later. The present article shows how to use the matching circles to explicitly reconstruct the global topology of space. The microwave background is isotropic to 1 part in $10^5$~\\cite{bennet}, which implies that the curvature of space is constant to 1 part in $10^4$~\\cite{ratra}. Our methods work equally well in the spherical, Euclidean, and hyperbolic cases. Current evidence suggests space is hyperbolic with $\\Omega_0$ approximately 0.3 or 0.4~\\cite{david}. If $\\Omega_0$ is 0.4, the LSS will have a radius of about 2 and enclose a volume of about 75, in units of the curvature radius. (The curvature radius provides a natural length scale in hyperbolic as well as spherical geometry. In spherical geometry it's usually called a ``radian'', so we will apply that term in the hyperbolic case as well.). Thousands of closed hyperbolic 3-manifolds of volume less than 7 are known~\\cite{jeff1,jeff2}; each would correspond to a universe in which the LSS encloses 10 or more images of each object in space, and in which the topology would be easily detectable. Moreover, the volume of a hyperbolic 3-manifold is a good measure of its complexity, and a least action argument~\\cite{gary,cgw} suggests that low-volume universes are more probable than high-volume ones. We don't rely on such arguments, but they give us hope that the cosmic topology will be detectable. ", + "conclusions": "" + }, + "9802/astro-ph9802297_arXiv.txt": { + "abstract": "Gravitational microlensing of quasars by stars in external galaxies can introduce fluctuations in the centroid of the ``point-like'' macro--images. The induced shifts are extremely small, on micro--arcsecond scales, below the limits of current optical observations. However, such shifts will become measurable with the proposed ``Space Interferometry'' mission, due to fly in 2005. The degree of the centroid shifts and their application as probes of both quasar structure and the stellar mass function in the lensing galaxy are discussed. ", + "introduction": "Over the last decade gravitational microlensing has shown itself to be not only a probe of the mass structure in our galactic halo (Alcock et al. 1993), but also of the scales of structure in active galaxies (e.g. Saust 1994, Lewis et al. 1997). Statistical studies of numerical simulations have revealed that long term light curves of microlensed quasars will additionally probe the mass function of the microlensing objects (Lewis \\& Irwin 1996). In this paper, a complementary probe of these properties, via the identification of microlensing induced ``centroid-shifting'' in the images of macrolensed quasars, is presented. The paper begins with a description of gravitational microlensing due to stars in external galaxies (Section~\\ref{microlensing_section}). The microlensing model is described in detail in Section~\\ref{microlensingmodel_section}, with the results of this study being presented in Section~\\ref{results_section}. These are discussed in the context of space-borne interferometry in Section~\\ref{space_section} and the conclusions of this study are presented in the final section. ", + "conclusions": "This paper has investigated the r\\^{o}le of microlensing in inducing image shifts in the centroids of the star--like images of multiply imaged quasars. It was found that for the quadruple quasar system, Q2237+0305, that substantial shift of $\\sim 100\\mu\\scmd$ are possible on time scales of months. Although not observable with current telescopes, such shifts will be detectable in the next decade with the proposed Space Interferometry Mission. This initial study has revealed that the characteristics of centroid shifting are sensitive to the size of the source undergoing microlensing, providing an independent probe of the scales of structure in the continuum emitting region of quasars. Similarly, these characteristics should depend on the mass function of the microlensing objects, although a proper determination of any influence will require statistical study of an ensemble of simulations. We wish to emphasize, however, that the light curves and image shifts presented here are drawn from only single realizations of the microlensing in the images of Q2237+0305. The results must, therefore, be taken as purely illustrative of the effect. Currently, an ensemble of simulations are underway; once complete, these will be employed in a study of the statistical properties of image centroid shifting and their dependence upon the scale of the source and mass function of the microlensing bodies." + }, + "9802/astro-ph9802318_arXiv.txt": { + "abstract": "We have measured the angular correlation function \\wth\\ for a sample of 871 Lyman-break galaxies (LBGs) in five fields at redshift $z\\sim 3$. Fitting the power-law $A_w\\theta^{-\\beta}$ to a weighted average of \\wth\\ from the five fields over the range $12\\simlt\\theta\\simlt 330$ arcsec, we find $A_w\\sim 2$ arcsec$^{\\beta}$ and $\\beta\\sim 0.9$. The slope is, within the errors, the same as for galaxy samples in the local and intermediate redshift universe, and a slope $\\beta=0.25$ or shallower is ruled out by the data at the 99.9\\% confidence level. Because $N(z)$ of LBGs is well determined from 376 spectroscopic LBG redshifts, the real-space correlation function can be accurately derived from the angular one through the Limber transform. The inversion of \\wth\\ is rather insensitive to the still relatively large uncertainties on $A_{\\omega}$ and $\\beta$, and the spatial correlation length is much more tightly constrained than either of these parameters. We estimate $r_0=3.3$\\er{0.6}{0.7} (2.1\\er{0.5}{0.4}) $\\hh$Mpc (comoving) for $q_0=0.1$ (0.5) at the median redshift of the survey, $\\bar z=3.04$ ($h$ is in units of 100 km s$^{-1}$ Mpc$^{-1}$ throughout this paper). The observed comoving correlation length of LBGs at $z\\sim 3$ is comparable to that of present-day spiral galaxies and is only $\\sim 50$\\% smaller than that of present-day ellipticals; it is as large or larger than any measured in recent intermediate-redshift galaxy samples ($0.3 \\simlt z \\simlt 1$). By comparing the observed galaxy correlation length to that of the mass predicted from CDM theory, we estimate a linear bias for LBGs of $b\\sim 1.5$ ($4.5$) for $q_0=0.1$ ($0.5$), in broad agreement with our previous estimates based on preliminary spectroscopy. The strong clustering and the large bias of the LBGs are consistent with biased galaxy formation theories and provide additional evidence that these systems are associated with massive dark matter halos. The results of the clustering of LBGs at $z\\sim 3$ emphasize that apparent evolution in the clustering properties of galaxies may be due as much to variations in effective light-to-mass bias parameter among different galaxy samples as to evolution in the mass distribution through gravitational instability. ", + "introduction": "In most cosmological models, galaxies are expected to be biased tracers of the underlying mass-density field, with the level of light-to-mass bias being a function of the galaxy mass; more massive galaxies would tend to populate volumes of space with a higher overall mass density, and, as a consequence, would be characterized by a stronger spatial clustering than less massive systems (e.g. Kaiser 1984; Mo \\& White 1996). Furthermore, the bias of galaxies respect to the mass is expected to evolve with cosmic time as a result of gravitational growth of density perturbation and hierarchical merging (Matarrese \\et 1997; Mann, Peacock \\& Heavens 1997; Bagla 1997). Empirically, it has been known for some time that different types of galaxies do indeed cluster differently; numerous large galaxy redshift surveys (Davis \\et 1988; Hamilton 1988; Santiago \\& Da Costa 1990; Loveday \\et 1995; Tucker \\et 1996; Valotto \\& Lambdas 1997) in the local universe have shown that early-type galaxies (E/S0) are more strongly clustered than later types (Sp/Irr), with a two-point correlation function that is generally steeper and a correlation length $\\sim 2$ times larger. A similar trend with the absolute luminosities of the galaxies has also been observed (Park \\et 1994). In the past few years, several deep redshift surveys have probed field galaxies in the intermediate-redshift universe (e.g. Lilly \\et 1995; Cowie, Hu \\& Songaila 1995). These find similar clustering segregation as in the local universe, with the redder and more luminous systems more strongly clustered than their bluer and fainter counterparts (Le F\\`evre \\et 1996; Carlberg \\et 1997), and detect apparent evolution in galaxy clustering, with a comoving correlation length $r_0$ that is three times smaller at $z\\sim 1$ than in local samples. If the galaxies in these samples at different redshifts all had the same bias with respect to the mass distribution, then the observed differences in galaxy clustering trace the evolution of mass clustering, and could be used to constrain cosmology; however, it is difficult to understand the mix of galaxy masses included in magnitude limited surveys as a function of redshift. One might hope that a sample's bias would not depend strongly on how it was selected, but if this were the case then different redshift surveys would currently be in quantitative disagreement with each other (Carlberg \\et 1997). It seems likely, then, that a sample--dependent (both because of redshift effects and selection criteria) light-to-mass bias could be at least partly responsible for the observed ``evolution'' of galaxy clustering with redshift, and in this case it would be difficult to draw cosmological conclusions from those surveys. Recently, it has become possible to identify large numbers of galaxies in a narrow redshift range using photometric techniques (e.g. Connolly \\et 1995, Steidel \\et 1996a, Madau \\et 1996). In contrast to traditional magnitude--limited surveys, which contain a wide range of galaxies over a large interval of time, and likely different mixtures of galaxies at different redshifts, a sample selected in this way provides a snapshot of the locations of similar galaxies over a small span of time. As a result, the observed clustering is much easier to interpret. An example of a photometric redshift technique is the Lyman-break technique (Steidel \\& Hamilton 1993, Steidel, Pettini, \\& Hamilton 1995, Steidel \\et 1996a, Giavalisco \\et 1996, Madau \\et 1996) which selects the brightest star-forming (and relatively dust-free) galaxies at high redshift. It is still unclear what the lower redshift counterparts of these Lyman-break galaxies would be, and so one cannot easily draw cosmological conclusions by comparing the clustering strength of Lyman-break galaxies to the clustering strength of a lower-redshift sample; but we can use the sample for the more modest goal of constraining theories of galaxy formation. In particular there is a great deal of fruitful work to be done comparing the properties of this well-defined high-redshift sample with the predictions of numerical and semi-analytic models (Baugh \\et 1998, Jing \\& Suto 1998, Governato \\et 1998). In a previous paper (Steidel \\et 1998, Paper 1) we described a large concentration of LBGs in redshift space discovered in one of our survey fields, and argued that such a concentration would not exist in standard CDM cosmogonies unless LBGs were very biased tracers of mass. In the present paper, we present a complementary angular clustering analysis of the LBG candidates in 5 of our survey fields, which can be used in conjunction with the spectroscopic redshift distribution to estimate the spatial correlation function of the Lyman-break population at $z \\sim 3$. Again, we will find that these galaxies are much more strongly clustered than the mass would be according to models of hierarchical structure formation. Such a strong clustering of forming galaxies is actually in agreement with predictions of models in which LBGs are associated with relatively rare and massive dark matter halos (e.g., Baugh \\et 1997, Mo \\& Fukugita 1996, Jing \\& Suto 1998). ", + "conclusions": "We have measured the angular correlation function \\wth\\ of Lyman-break galaxies at redshift $z\\sim 3$. Fitting the power-law $A_w\\theta^{-\\beta}$ to a weighted average of \\wth\\ from the five fields over the range $12\\simlt\\theta\\simlt 330$ arcsec, we find $A_w\\sim 2$ arcsec$^{\\beta}$ and $\\beta\\sim 0.9$. The slope is, within the errors, the same as for galaxy samples in the local and intermediate redshift universe, and a slope $\\beta=0.25$ or shallower is ruled out by the data at the 99.9\\% confidence level. Because the redshift distribution $N(z)$ of LBGs is well determined from 376 spectroscopic redshifts, we have derived the real-space correlation function from the angular one through the Limber transform. The inversion of the \\wth\\ is rather insensitive to the still relatively large uncertainties on $A_{\\omega}$ and $\\beta$, and the spatial correlation length $r_0$ is much more tightly constrained than these. Using Monte Carlo simulations to derive the $1\\sigma$ error bars from the 68\\% confidence interval, we estimate $r_0=3.3$\\er{0.6}{0.7} (2.1\\er{0.5}{0.4}) $\\hh$Mpc (comoving) for $q_0=0.1$ (0.5) at the median redshift of the survey, $\\bar z=3.04$. Thus, the observed comoving correlation length of LBGs at $z\\sim 3$ is comparable to that of present-day spiral galaxies and is only $\\sim 50$\\% smaller than that of present-day ellipticals; it is as large or larger than any measured in recent intermediate-redshift galaxy samples ($0.3 \\simlt z \\simlt 1$). By comparing the observed correlation length of LBGs to that of the mass predicted from CDM theory, we have estimated a linear bias for LBGs of $b\\sim 1.5$ ($4.5$) for $q_0=0.1$ ($0.5$), in broad agreement with our previous estimates based on preliminary spectroscopy (Paper 1). The strong clustering and the large inferred bias of the LBGs are consistent with biased galaxy formation theories and provide additional evidence that these systems are associated with massive dark matter halos. The evolution of the slope of the correlation function of the mass (or, equivalently, that of the power spectrum at small scales) has a pronounced dependence on $\\Omega$. If the biasing parameter is a weak function of the spatial scale, the measured slope of the correlation function of LBGs, $\\gamma=1.98$\\er{0.28}{0.32}, is inconsistent with the predictions of the standard CDM theory with $\\Omega=1$ at the 99.9\\% confidence level. N-body simulations seem to suggest that the bias will not be strongly scale-dependent; however, until more is known about the scale-dependence of the bias, this conclusion will remain speculative. The results of the clustering of LBGs at $z\\sim 3$ emphasize that apparent evolution in the clustering properties of galaxies may be due as much to variations in effective light-to-mass bias parameter among different galaxy samples as to evolution in the mass distribution through gravitational instability. Our study shows that the traditional power-law model $\\xi(r,z)=\\xi_0(r)\\times (1+z)^{-(3+\\epsilon)}$ traditionally used to describe the gravitational evolution of clustering (Peebles 1980) is not a good representation for any value of $\\epsilon$ over the redshift interval $0\\simlt z\\simlt 3$. \\vskip1cm We are grateful to an anonymous referee for useful comments and suggestion that have improved our paper. We would like to thank the staff at the Palomar, Kitt Peak and Keck observatories for their invaluable help in obtaining the data that made this work possible. We are greatly indebted to Ed Carder at KPNO for understanding the cause of the $CuSO_4$ leakage in the $U_n$ filter and his superb work in repairing it. We have benefited from stimulating conversations with Ray Carlberg, John Peacock and Cedric Lacey, who have also kindly made available to us their theoretical models of CDM clustering evolution. We are grateful to Richard Ellis for his useful comments on an early version of the paper. MG has been supported by the Hubble Fellowship program through grant HF-01071.01-94A awarded by the Space Telescope Science Institution, which is operated by the Association of Universities for Research in Astronomy, Inc. under NASA contract NAS 5-26555. CCS acknowledges support from the U.S. National Science Foundation through grant AST 94-57446, and from the Alfred P. Sloan Foundation." + }, + "9802/astro-ph9802175_arXiv.txt": { + "abstract": "We present {\\it ROSAT} PSPC observations of the twin-jet radio galaxy 3C\\,449. The soft X-ray emission from this object is dominated by an extended halo with a scale comparable to that of the radio source. The asymmetry of the X-ray emission is reflected in that of the radio lobes, providing evidence that the behaviour of the jets is strongly influenced by the external medium. A region of reduced X-ray surface brightness coincident with the southern radio lobe of 3C\\,449 suggests that the radio source has displaced thermal plasma from the X-ray-emitting halo. However, the minimum pressure in the radio lobe is considerably lower than our estimates of the pressure in the external medium. We discuss the implications for the dynamics of the radio source. ", + "introduction": "3C\\,449 (B2 2229+39) is a well-known FRI radio galaxy (Fanaroff \\& Riley 1974) at a redshift of 0.0171. Its symmetrical inner jets have been well studied in the radio (Perley, Willis \\& Scott 1979; Cornwell \\& Perley 1984). On larger scales, the southern jet flares into a lobe, while the northern jet continues to be collimated until it fades into the noise (Birkinshaw, Laing \\& Peacock 1981; J\\\"agers 1987; Andernach \\etal\\ 1992; Leahy, Bridle \\& Strom 1997). 3C\\,449's host galaxy, UGC 12064, has a close bright companion (e.g.\\ Balcells \\etal\\ 1992) and models have been proposed in which the orbital motion of the host is responsible for the symmetrical oscillations in the jets (e.g.\\ Lupton \\& Gott 1982; Hardee, Cooper \\& Clarke 1994). UGC 12064 appears in the optical to lie in a poor cluster (Zw 2231.2+3732), and the extended optical halo surrounding it and its companion allows it to be classed as a cD galaxy (Wyndham 1966). Miley \\etal\\ (1983) used the {\\it Einstein} IPC to show that it is a weak, extended X-ray source, and the rotation measure gradients in the jet are thought to be produced by the hot gas responsible for the X-ray emission (Cornwell \\& Perley 1984). In this paper we present new, sensitive observations of 3C\\,449 with the {\\it ROSAT} PSPC, which allow us to map in detail the hot gas halo surrounding the source. Throughout we use a cosmology in which $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$, $q_0 = 0$. At the distance of 3C\\,449, 1 arcsec corresponds to 0.485 kpc. ", + "conclusions": "We have presented evidence that the large-scale structure of 3C\\,449 is being determined by the distribution of the hot intergalactic plasma. The radio lobes are significantly underpressured if the radio-emitting plasma is electron/positron with filling factor unity and the magnetic field and particles are in energy equipartition. It is then hard to understand how the southern lobe can be associated with a deficit of X-ray emission. This problem, which seems to be common in low-power radio sources, may be resolved in a number of ways; the most plausible seem to be a dominant contribution to the particle energy density from relativistic or thermal (entrained) protons or a plasma filling factor much less than unity." + }, + "9802/astro-ph9802343_arXiv.txt": { + "abstract": "\\parindent=0.6cm We examine whether a cosmologically significant distribution of dark galaxy group or cluster-sized objects can have an optical depth for multiple imaging of distant background sources which is comparable to that from known galaxies while at the same time producing angular splittings of the same order of magnitude. Our purpose is to explore whether such objects could realistically account for some of the observed lenses. Modeling such systems as isothermal spheres with core radii, and assuming a Schechter-type distribution function, we find that independent of the cosmology (open, flat matter dominated, or flat cosmological constant dominated) an allowed parameter range exists which is comparable in velocity dispersion to that for known compact groups of galaxies, although the preferred core radii are somewhat smaller than that normally assumed for compact groups. Dark cluster-sized objects, on the other hand, cannot reproduce the observed lensing characteristics. If the one known Dark cluster were a good representative of such a distribution, most such objects would not produce multiple images. We also present a result for the angular splitting due to an isothermal sphere lens with non-zero core radius, extending earlier work of Hinshaw and Krauss (1987). Our results are expressed as contour plots for fixed lensing probabilities, and angular splittings. ", + "introduction": "\\ec Ever since the first multiply imaged quasar was observed, it was clear that the statistics of gravitational lensing could be utilized to pin down cosmological parameters (i.e. \\cite{TOG}; \\cite{Ho}). Specifically, for a given cosmological model, one can predict the optical depth due to lensing by normal galaxies (presuming one has a model of normal galaxies and their distribution), and compare that with observations. Several large scale surveys have been performed, searching for multiple imaging of distant quasars by intervening galaxies, and more than a dozen such lensing events have been observed (\\cite{kochanek96}). However, there is one slight peculiarity. While the overall frequency of lensing events, and the rough angular splittings are reasonably consistent with model expectations, in a significant fraction of the cases, the actual lensing galaxy is not visible. Given the predicted mass of the lensing systems, it is not obvious that such a large fraction should remain unresolved. This prompts the natural question: Could the lenses be dark objects, perhaps objects related to the distribution of dark matter in the Universe---perhaps failed galaxies? This is not a new idea, and it is one which has been beset with problems. In the first place, anything close to a closure density of compact objects generally produces an optical depth for lensing which is too large (\\cite{pressgunn}). Second, if the dark objects are failed galaxies, then their properties will generally preclude producing multiple images with the observed angular splittings (\\cite{HK}). Here we examine another possibility. Could larger systems, on the scale of groups or clusters account for an observable fraction of the known lenses? The recent observation of a cluster-sized mass distribution containing one luminous galaxy (\\cite{nature}; \\cite{Mushotzky}) provides some additional {\\it a posteriori} motivation for considering this hypothesis. On first glance such a possibility seems implausible. Larger systems, with larger velocity dispersions, will produce larger angular splittings, if they produce multiple images at all. Thus, it seems unlikely that such systems might reproduce the observed lensing characteristics, which, as alluded to above, are comparable to those one would predict for the known population of galaxies. However, Hinshaw and Krauss (1987; hereafter HK) demonstrated that under certain conditions a finite core radius suppresses the mean angular splitting due to isothermal sphere mass distributions. Here, we generalize the earlier HK result, and prove that this approximation is good for all lensing impact parameters inside the critical disk for multiple image formation. We then demonstrate that for a reasonable range of velocity dispersions, core radii, and total mass, assuming a Schechter-type distribution function, such systems can produce optical depths for lensing comparable to that due to known galaxies, with comparable angular splittings. It is well known that due to a variety of selection effects (magnification biasing, etc.) the actual fraction of strong lensing events in any sample can differ dramatically from the naive optical depth calculation. However, because these systematic effects should be largely independent of the nature of the lenses themselves, if they are producing comparable angular splittings, etc, we need not consider these effects in detail here. In particular, we compare the calculated optical depth for lensing by dark objects to the naive optical depth for lensing by known galaxy distributions, known to be in the range of $10^{-3}$ to $10^{-4}$. Presumably, if the other selection effects are comparable, this will then imply that such systems could produce at least some fraction of the observed events. We thus derive contour plots in the parameter space of velocity dispersion and redshift for fixed total optical depth in order to explore the suggested range of dark lenses. For this range we then explore the magnitude of the predicted mean angular splittings. We conclude with a brief summary of our results. \\bc ", + "conclusions": "\\ec Over a dozen multiply imaged quasars have now been observed in various optical and radio surveys. While these have angular splittings characteristic of those one might expect to be induced by galaxies (i.e. $O(2''$-$3'')$), several have no observable lensing systems, even when one might expect the lensing galaxy to be resolvable. When one calculates the optical depth for lensing by known galaxy distributions, ignoring magnification biasing, one typically finds $ 10^{-4} \\le \\tau \\le 10^{-3}$. Thus, one might expect that a predicted optical depth of as low as $10^{-5}$ for some other distribution would produce at least one or two lenses in the existing surveys. In fitting the optical depth, however, one must confront the tension between optical depth and angular splitting discussed above (also compare Fig.~1 and Fig.~4). Systems with velocity dispersion greater than about $450 km/sec$ would require the expectation value of the angular splitting to be less than half that predicted for an equivalent singular isothermal sphere. This, in turn, requires $\\rc$ to be larger relative to $\\sigma$, which in turn, however, suppresses the optical depth. Nevertheless, in Table 1 we display several sample values of core radii and velocity dispersions which would be expected to produce a fraction of the observed lenses with angular splittings comparable to those observed. In this table we have fixed $\\rr =10$, set the source redshift to 3, and fixed the parameters so that the mean predicted angular splitting is $2.5$ arcsec. Note that the parameter range is comparable to compact groups of galaxies, which have compatible $\\frac{(\\sigma/km/sec)^{2}}{\\rc/(h^{-1}Mpc)}$. For example, one study suggests $\\sigma$ is around 331 $km/sec$, and $\\rc$ is around $15 h^{-1} kpc$ (\\cite{mog}; \\cite{PBE}). These numbers give $\\frac{(\\sigma/km/sec)^{2}}{\\rc/(h^{-1}Mpc)} \\approx 7.3*10^{6}$, as compared to the preferred values of $\\frac{(\\sigma/km/sec)^{2}}{\\rc/(h^{-1}Mpc)}$ between $10^{7}$ and $10^{9}$ to result in an optical depth comparable to that of the known galaxy distribution. Also note that if we were to integrate over a distribution of such objects, our earlier arguments suggest that comparable optical depths and splittings could be obtained even if the mean value of $r_c$ were somewhat larger than given in the table. However cluster-sized objects do not seem to be viable candidates for dark lenses because of the larger predicted splittings when multiple image formation does occur, and more importantly because the predicted optical depth is too small. (The $\\rc$ of a regular cluster is large, but the velocity dispersion is only about 2.5 times larger than that of compact groups of galaxies (\\cite{zghv}). ) This conclusion is reinforced by the recent observation of an actual dark cluster by Hattori et. al. (1997) in the lensing system, MG2016+112. Based on the inferred mass of the object using X-Ray estimates of the potential, its size, and the size of the core radius of this system, we find that $\\beta_0 > 1/2$, implying that the cluster is not responsible for the observed multiple images in this system. This is again supported by the small angular splitting of $3.4$ arcsec between the images, which is characteristic of the one observed galaxy in this system. Thus, perhaps paradoxically, the only known example of a large-scale dark object suggests that such cluster-scale objects, even if they have a significant mass density in the universe, are probably largely irrelevant for the statistics of multiply imaging distant quasars. Rather, dark objects on intermediate scales, between galaxy and cluster scales, are more likely possible candidates for dark lenses. A very recent paper which has shown a very high rate of galaxy lensing in radio surveys may be significant in this regard (\\cite{Helbig}). We also note both that a recent study of the luminosity function of the compact groups of galaxies also lends support to our choice of Schechter $\\alpha_{1} = -1$ (\\cite{ZCR}), and that the currently favored cosmological model involving a flat universe with cosmological constant (i.e. \\cite{krauss}) produces the largest optical depths at high redshift for given $\\frac{(\\sigma/km/sec)^{2}}{\\rc/(h^{-1}Mpc)}$, as one would expect, based on the increase in the optical depth for lensing by galaxies in this cosmology. Clearly, in order to know whether dark clusters are important for lensing in the actual universe, larger surveys will be required, in order to reliably determine how many lensing events might not be associated with galaxies. If such events continue to be observed, our analysis suggests that the distribution of angular splittings will be an important observable which might constrain possible models. \\bc {\\bl ACKNOWLEDGMENTS} \\ec" + }, + "9802/astro-ph9802033_arXiv.txt": { + "abstract": "Evolution in the merger rate as a function of redshift is {\\em in principle} the key observable testing hierarchical models for the formation and evolution of galaxies. However, {\\em in practice}, direct measurement of this quantity has proven difficult. In this opening review I outline the current best estimates for the merger rate as a function of cosmic epoch, focussing mostly upon recent advances made possible by deep ground-based redshift surveys and morphological studies undertaken with HST. I argue that a marriage of these techniques, in an attempt to determine the space density of mergers amongst the abundant morphologically peculiar population at high redshifts, is probably the most promising currently-available avenue for determining the prevalence of mergers at high redshifts. However, resolved kinematical studies, which seem set to become available in the next few years, are probably the best hope for a definitive determination of the space density of mergers at high redshifts. ", + "introduction": "Three observational techniques have been used to probe changes in the merger rate as a function of cosmic epoch. These are (a) studies of angular and physical correlation functions, (b) pair counts, and (c) morphological studies. The advantages ($+$) and disadvantages ($-$) of each of these approaches for studying high-redshift mergers can be crudely summarized as follows. \\subsection{Correlation Functions } \\begin{list}{$+$}{} \\item A close connection to large-scale structure work via the clustering statistics $w(\\theta)$ and $\\xi(r)$. Since mergers can be identified as the end-products of large-scale clustering, changes in the correlation function at small radius provide a conceptual link between the large-scale and small-scale regimes being probed with the same statistics. \\item Biases inherent in measuring correlation functions are fairly well-understood. \\end{list} \\begin{list}{$-$}{} \\item Measurement of correlation functions is best suited to large samples. \\item These statistics are hard to measure at small radii, particularly when galactic components (such as giant HII regions) become difficult to distinguish from merging companions. This is often the case when probing distant galaxies in the rest-frame ultraviolet, {\\it eg.} in the Hubble Deep Field, where knots of star-formation cannot easily be distinguished from companions (Colley et al. 1996). \\item {\\em Not really measuring merger rate (see below)} \\end{list} \\subsection{Pair Counts Approach} \\begin{list}{$+$}{} \\item Simple statistics to measure (described in detail in the next section). \\item Conceptually the pair-counts approach is an integration over the two-point correlation function, with better signal-to-noise properties at small radii, so the biases are also fairly well-understood. \\end{list} \\begin{list}{$-$}{} \\item Like correlation functions, the statistic becomes ambiguous when merging companions become indistinguishable from galactic components. \\item {\\em Not really measuring merger rate (see below)} \\end{list} \\hspace{0.5cm} \\noindent The most important criticism common to both the correlation function and pair counts approaches is that neither provides a direct measure of the merger rate. The fundamental difficulty is {\\em the uncertain physical timescale over which merging occurs}, given physical proximity between galaxies. Both statistics probe the fuel ``reservoir'' available for close gravitational interaction, but what we really seek is an understanding of the rate at which the merger ``engine'' operates in converting these galaxies into the by-products of mergers. To understand this, one must observe the mergers in progress. So the third approach must necessarily be a morphological one: \\subsection{Morphological Approach} \\begin{list}{$+$}{} \\item Direct observation of {\\em mergers in progress} \\end{list} \\begin{list}{$-$}{} \\item Poorly-understood biases ({\\em eg.} morphological K-corrections) \\end{list} Because of space limitations, and because they are well-reviewed elsewhere, this review ignores studies of the correlation function (eg. Neuschaefer et al. 1997), and touches only rather superficially on recent studies on pair-counts in \\S2, in order to focus mostly on summarizing very recent progress made on morphological studies at high-redshift with HST in \\S3. ", + "conclusions": "The most recent pair-count analyses seem to suggest that the fraction of physical pairs grows as $\\sim (1+z)^3$. However, the conversion between pair fraction and merger rate is uncertain, and the pair count work so far published is limited to fairly low redshifts ($\\sim0.3$). At higher redshifts morphological work with HST indicates that by $I=24$ mag something over 30\\% of all galaxies are morphologically peculiar. Simulations and follow-up spectroscopic work suggest this excess in morphologically peculiar systems is a physical effect, and not merely the result of ``morphological K-corrections''. The fraction of mergers amongst these morphologically peculiar galaxies is unknown, because obviously merging local systems, such as nearby major starburst galaxies, no longer appear like conventional mergers at high redshifts. A preliminary analysis of image bimodality (a robust parameter that in principle flags major mergers even at high redshift) amongst $I<22$ mag peculiar systems with $z<1$ suggests that around 40\\% the morphologically peculiar galaxies are strongly bimodal, and thus probably merging. This is consistent with a merger rate increases of $\\sim (1+z)^3$. Internal color analyses of morphologically peculiar systems is an interesting next step in probing high-redshift mergers, leading ultimately to kinematical investigations in the next few years." + }, + "9802/astro-ph9802205_arXiv.txt": { + "abstract": "Microlensing is a powerful tool for studying stellar atmospheres because as the source crosses regions of formally infinite magnification (caustics) the surface of the star is resolved, thereby allowing one to measure the radial intensity profile, both photometrically and spectroscopically. However, caustic crossing events are relatively rare, and monitoring them requires intensive application of telescope resources. It is therefore essential that the observational parameters needed to accurately measure the intensity profile are quantified. We calculate the expected errors in the recovered radial intensity profile as a function of the unlensed flux, source radius, spatial resolution the recovered intensity profile, and caustic crossing time for the two principle types of caustics: point-mass and binary lenses. We demonstrate that for both cases there exist simple scaling relations between these parameters and the resultant errors. We find that the error as a function of the spatial resolution of the recovered profile, parameterized by the number of radial bins, increases as $N_R^{3/2}$, considerably faster than the naive $N_R^{1/2}$ expectation. Finally, we discuss the relative advantages of binary caustic-crossing events and point-lens events. Binary events are more common, easier to plan for, and provide more homogeneous information about the stellar atmosphere. However, a sub-class of point-mass events with low impact parameters can provide dramatically more information provided that they can be recognized in time to initiate observations. ", + "introduction": "Originally proposed by Paczy\\'nski (1986) as a method to detect the presence of massive compact objects in the halo of our Galaxy, microlensing has increasingly been recognized as a tool for studying a broad range of astrophysical phenomena. Various applications include detection and characterization of binary and planetary systems (Mao \\& Paczy\\'nski 1991; Gould \\& Loeb 1992), reconstruction of the stellar mass function down to masses below the hydrogen burning limit (Paczy\\'nski 1991; Griest et al.\\ 1991; Gould 1996), measurement of the rotation speed of giants (Gould 1997), measuring the transverse velocity of galaxies (Gould 1995a), and probing the central engines of quasars (Gould \\& Gaudi 1997). Recently, Valls-Gabaud (1994, 1997) and Sasselov (1996) have proposed using microlensing to study stellar atmospheres. Here we analyze this application in detail. Currently, three collaborations (MACHO, Alcock et al.\\ 1997a; OGLE, Udalski et al.\\ 1997; EROS, Ansari et al.\\ 1996) have ongoing projects that survey the Galactic bulge with roughly nightly sampling in order to detect microlensing events. Over 60 events per year are being detected. These data are being analyzed real time which has allowed MACHO to issue ``alerts'', notifications of ongoing events detected before the peak. Two follow-up collaborations (PLANET, Albrow et al.\\ 1997, 1998a; GMAN, Alcock et al.\\ 1997b) have formed in order to monitor these alerts around the clock with high precision and high temporal sampling with the aim of detecting (and further alerting on) deviations from the standard microlensing light curve, such as would be expected from binary lenses, binary sources, finite sources, and parallax. These deviations are useful in that they can provide additional information about the lens and/or source. In addition, a collaboration has formed to conduct {\\it spectroscopic} monitoring of alerted events in order to study the source stars in detail and has observed several events to date (Lennon et al.\\ 1997). Thus, the prospect for the real-time detection and monitoring of light curve anomalies both photometrically and spectroscopically is promising. All gravitational microlenses have caustics, defined as the set of points in the source plane where the magnification of a point source is formally infinite. When a finite source crosses a caustic, the gradient of the magnification over the source is large, and therefore different parts of the source are magnified by different amounts. Hence, the source is partially resolved. Different parts of the source are resolved at different times during the caustic crossing, and thus, by taking a series of measurements during the course of the crossing, one can recover the intensity profile of the source. Several workers realized that finite source effects could be useful for breaking or partially breaking the degeneracy among microlensing event parameters (Gould 1994; Nemiroff \\& Wickramasinghe 1994; Witt \\& Mao 1994; Maoz \\& Gould 1994) and that variations in the surface profile could be exploited to this end (Witt 1995; Loeb \\& Sasselov 1995; Gould \\& Welch 1996). However, Valls-Gabaud (1994, 1997), Sasselov (1996), and Heyrovsky, Sasselov, \\& Loeb (1998) proposed to exploit the same effects for radically different aims. Instead of using finite-source effects to learn more about the lens, they sought to learn more about the source. The basic idea is as follows: Imagine that one could image separately different annular rings on the surface of a star. In effect one would be sampling different depths of the photosphere. Since the temperature varies as a function of depth, the broad spectral energy distribution would change with annulus, with more blue light near the center (greater depths) and more red light near the outer limb (lesser depths). That is, the star would be limb-darkened, and more in the blue than the red. Since different spectral lines form at different depths, one would expect that the detailed spectral profile would vary as a function of annular radius. Hence, the entire atmosphere could be studied as a function of depth by resolving the two-dimensional (radius and wavelength) spectral profile of the star. Currently, it is only possible to study stellar atmospheres in this way for the Sun and eclipsing binaries. However, since the surface of the source star is partially resolved during a microlensing caustic crossing, one can also probe the atmospheres of the source stars for these types of events. Both Valls-Gabaud (1994, 1997) and Sasselov (1996) used specific stellar atmosphere models to construct broad-band and spectral line brightness profiles, and then used these profiles to predict in detail the variations in the broad-band color or equivalent width of specific lines that one would expect during the course of a point-mass microlensing event. Sasselov (1996), and Heyrovsky et al.\\ (1997) also consider the effects of star spots on the microlensing light curve. In addition, they compared their predictions to MACHO Alert event 95-30, a point-mass caustic crossing event for which spectra were taken during the course of the crossing, and for which variations in the optical TiO bands were detected (Alcock et al.\\ 1997b). All of these authors predict that the color and spectral-line variations during the caustic crossing should be significant and note that this provides an entirely new method of studying stellar atmospheres. Although caustic crossing events are in principle useful for studying the atmospheres of stars, these events are rare, and they typically last for only about 7 hours (for a giant source). For this method to be successful, it is essential that observers have a clear sense of what can be accomplished with these events, since substantial telescope resources are likely to be expended. To this end, we approach this topic from another perspective. We quantify the intrinsic ability of both point-mass and binary lens microlensing caustic crossings to recover the radial variation of the intensity of the star for any arbitrary wavelength, and hence, for any spectral line. Specifically, we calculate the fractional error in the recovered intensity profile as a function of the unlensed flux of the source, the duration of the measurements, the size of the telescope, the magnification of the event, and the spatial resolution of the recovered intensity profile. This information will be useful to observers in making rational decisions about which events to follow, and what resources are required to address specific questions. ", + "conclusions": "While both point-mass lenses and binary lenses can in principle be used to resolve the two-dimensional (radial + spectral) profile of a star, binary lenses are substantially easier to use. First, for a binary one always has warning of the second caustic crossing. When the source crosses the caustic the first time it is suddenly magnified by a factor $\\rho^{-1/2}\\sim 7$ and hence is easily recognized. The second crossing can then be expected in several days to several weeks. Intensive photometric monitoring (now routinely undertaken by PLANET and GMAN) can then be analyzed to make a more precise prediction. From Figure 4b it is clear that the most useful portion of the second caustic crossing is the final $\\sim 70\\%$ of the time that the source actually straddles the caustic. The onset of this optimal period can be judged extremely accurately if photometric monitoring is proceeding simultaneously, and reasonably accurately even one day in advance. By contrast, there is no way to guarantee a priori that a point-mass caustic crossing will occur because one does not know the size of the Einstein ring projected onto the source plane beforehand, and hence one does not know $\\rho$. Using optical photometry alone, one can ``predict'' a source crossing only at about the time it begins. Using optical/infrared photometry, it could be predicted at $r\\sim 1.5\\,\\rho$ (Gould \\& Welch 1995), but this would leave only a few hours' warning. Second, fold caustics generically provide information about the entire radial profile of the star while point caustics provide information only for annuli of the star that are greater than the impact parameter, $\\beta$. See Figure 4. Again, it is virtually impossible to predict in advance for which events $\\beta\\ll \\rho$ (and hence for which events one can resolve essentially the whole star), although once the peak occurred, these events could be recognized provided that the star was being monitored photometrically. In any event, of all point-caustic transits, only a fraction $\\beta/\\rho$ have impact parameters smaller than $\\beta$. Third, binary-lens events are probably more common than point-mass caustic crossings. The fraction of events with binary-caustic crossings has not yet been established empirically, but 5\\% appears to be a plausible estimate. The fraction of point-mass caustic crossings is $\\langle \\rho\\rangle$, but by the argument of the previous paragraph, only about half of these are really useful. The mean radius of a giant is $R\\sim 22\\,R_\\odot$ (Gould 1995b), about 2.2 times larger than the fiducial value used in equation (4.1). Thus, the fraction of events with useful point-mass caustic crossings is $\\sim 0.5\\langle \\rho\\rangle\\sim 2\\%$. On the other hand, as shown by Figure 4, over the range probed by the point lens ($r<\\beta$), the point-lens errors are less than half as large as those of the binary. This advantage diminishes as $\\rho^{-1/4}$ for larger stars, but is still substantial for most giants. The problem of recognizing events with $\\beta\\ll\\rho$ sufficiently early to permit spectroscopic monitoring is formidable. However, if they can be monitored beginning at their peak, they are the best events to follow." + }, + "9802/astro-ph9802049_arXiv.txt": { + "abstract": "The astro\\-physically very interesting \\lii~671\\,nm line has been observed with high spatial resolution in solar granulation with the intention to diagnose departures from local thermodynamic equilibrium (LTE) in the line formation. The spectral feature is very weak, so this is also a test on the limits of such observations. The observations of the \\lii\\ line and other weak lines nearby are compared with the results of synthetic line calculations in three-dimensional granulation simulations. The dependence of line-centre velocities on photospheric continuum brightness is well described by the simulations. The observed equivalent width of the \\lii\\ line show an approximately flat dependence on continuum brightness, contrary to the theoretical LTE results. Detailed modelling of the line radiative transfer, with an approximate inclusion of three-dimensional effects, gives a better agreement with observations. The match is not perfect and various interesting reasons for this are considered. However, the possibility of systematic errors caused by the sensitivity of the \\lii\\ equivalent width to continuum placement calls for cautiousness in the conclusions. ", + "introduction": "Stellar lithium abundances are the subject of many current debates in astrophysics (e.g., Thorburn 1996, Crane 1995, Spite \\& Pallavicini 1995). This paper continues the study of \\lii\\ line formation reported by Kiselman (1997), hereafter referred to as Paper\\,I. The aim is to improve the understanding of lithium abundance determinations of solar-type stars by investigating the formation of lithium lines outside the classical realm of plane-parallel homogeneous photospheres and local thermodynamic equilibrium (\\lte). It is also intended to test the idea put forward in Paper\\,I that spatially resolved solar spectroscopy can be used to test \\nlte\\ results in less model-dependent ways than when just line profiles in integrated light are used. At the same time it tests the limits of such observational work since the lines studied are very weak. Paper\\,I was motivated by the pro\\-posal of Kurucz (1995) that lithium abundances derived for extremely metal-poor solar-type stars using 1D photospheric models and assuming \\lte\\ may be underestimated by as much as 1\\,dex due to 3D \\nlte\\ effects. The idea was that the transparency of ``cold'' regions in metal-poor photospheres would cause all lithium there to be ionised by ultraviolet radiation from ``hot'' regions. This would lead to a very weak line in the resulting mean spectrum, since the contribution to the spectral line from the hot regions is already small. No such large effect of granulation on abundance determinations has yet been demonstrated by observations or detailed simulations. The standard behaviour of lines over the {\\em solar} granulation pattern seems to be that lines -- irrespective of the ionisation stage -- get stronger in bright granules and weaker in darker intergranular lanes (Steffen 1989, Kiselman 1994). Thus one can expect that the influence on abundance ratios, which are derived from line ratios, will not be excessively large (Holweger et al. 1990). The investigation of \\lii\\ line formation in a 3D model of solar granulation in Paper\\,I demonstrated that the Kurucz mechanism at least does not work in solar granulation. This is also the conclusion of Uitenbroek (1998) who used a different \\nlte\\ treatment but the same kind of granulation model. As regards metal-poor stars, the final verdict must of course wait until we know more about the inhomogeneity properties of their photospheres. This paper will only concern the solar \\lii\\ lines and is organised as follows. First the simulations are discussed in somewhat more detail than in Paper\\,I. Then the solar observations are described, and finally the observations are compared with the simulation results. \\begin{figure} \\ifnum\\doepsf=1\\hspace*{.2cm}\\epsfxsize=8cm\\epsffile{fig1.eps} \\else\\picplace{1cm}\\fi \\caption[]{Results from line transfer calculations in a granulation snapshot using different treatments and approximations: \\lii\\ equivalent widths as function of continuum intensity (left) and the ratio of these to the corresponding \\lte\\ values (right). Labels are explained in text. The mean values given are intensity-weighted means \\label{fig_icwsol1}} \\end{figure} ", + "conclusions": "The observed $I_c - W$ relation for the \\lii\\ 671\\,nm line is rather flat and falls in between the theoretical \\lte\\ and the \\nlte\\ relations, though closer to the latter. It seems that the results exclude \\lte\\ as a line-formation hypothesis, something that is of interest in principle, though of course not very surprising. The \\nlte\\ modelling of this paper is at best only marginally confirmed, however. It is the extreme weakness of the \\lii\\ line that makes definite conclusions difficult to make. This is both because the necessary binning makes it impossible to study the spread in the $I_c - W$ plot and because the equivalent widths are sensitive to errors in the continuum placement. Spectra with better spatial resolution and/or higher S/N would probably allow firmer conclusions if they could be acquired. The idea that spatially resolved spectroscopy may be helpful in revealing \\nlte\\ effects is still viable. It should be tested on somewhat stronger lines of other elements, even though their \\nlte\\ modelling is intrinsically more complicated and uncertain than for the relatively simple \\lii\\ case. A coming paper (Kiselman, in preparation) will discuss similar solar observations and possible departures from \\lte\\ of a large set of Fe lines of different strengths." + }, + "9802/astro-ph9802080_arXiv.txt": { + "abstract": "Most X-ray novae (aka soft X-ray transients) contain black hole primaries. In particular, the large mass functions measured for six X-ray novae directly clinch the argument (within general relativity) that they contain black holes. These firm dynamical results are discussed, and the urgent need to determine precise masses for black holes is stressed. The dynamical evidence for black holes is convincing but it is indirect. Now it appears that {\\it direct} evidence may be at hand. Three recent studies have revealed phenomena that very likely probe strong gravitational fields: (1) a comparison of the luminosities of black hole systems and neutron star systems has yielded compelling evidence for the existence of event horizons; (2) RXTE observations of fast, stable QPOs have probed the very inner accretion disks of two black holes; and (3) three different types of low energy spectra have been linked to different black-hole spin states (e.g. Kerr {\\it vs.} Schwarzschild). ", + "introduction": " ", + "conclusions": "" + }, + "9802/astro-ph9802339_arXiv.txt": { + "abstract": "Cyg~X-1 exhibits irregular X-ray variability on all measured timescales. The usually applied shot noise models describe the typical short-term behavior of this source by superposition of randomly occuring shots with a distribution of shot durations. We have reanalyzed EXOSAT ME observations of Cyg~X-1\\ using the more general Linear State Space Models (LSSMs). These models, which explicitly take the observation-noise into account, model the intrinsic system variability with an autoregressive (AR) process. Our fits show that an AR process of first order can reproduce the system variability of Cyg~X-1. A possible interpretation is again the superposition of individual shots, but with a single relaxation time $\\tau$. This parameter was found to be 0.19\\,s. ", + "introduction": " ", + "conclusions": "\\label{kapitel5} \\subsection{The LSSM[1] results --- a generalization of the shot noise approach}\\label{kapitel5.1} We have shown that the short-term variability of the EXOSAT ME lightcurves of Cyg~X-1\\ can be well described by a Linear State Space Model of first order, i.e.\\ the dynamics of the system can be modeled as an autoregressive process with one temporal parameter -- the relaxation timescale $\\tau$. We find $\\tau$ to be (0.19$\\pm$0.03)\\,s for the ME energy-range. The relaxator is not found to be correlated either with time nor with the luminosity of the source, indicating that the physical process producing the emitted radiation is very stable and suggesting that the short-term variability is independent of the mass accretion rate of the system. The LSSM analysis in the time domain allows us to estimate the fraction of the countrate that is due to observation-noise. When analyzing the periodogram in the frequency domain this white noise component can only be represented by a constant whose extraction is not trivial (\\cite{belloni:90b,zhangw:95a}). Furthermore, the problem of spectral leakage that has to be delt with in the frequency domain, is circumvented by fitting in the time domain. For these reasons our LSSM analysis delivers results with higher statistical significance than corresponding frequency domain fits (\\cite{koenig:97a}). These LSSM[1] results allow a much simpler description of hard state short-term variability than multi-timescale shot noise models. Although a detailed quantitative comparison is beyond the scope of this paper, the reproduction of the periodogram by the LSSM[1] (compared to its approximation by adding shot-profiles with different timescales, corrected for observation-noise and binning) as well as the greater sensitivity of the LSSM fitting procedure (compared to frequency domain fits, which are usually used to evaluate shot noise models) suggest that LSSM[1]s are better suited to describe the nature of the observed variability. The LSSMs can model the different realizations of a stochastic relaxator $\\tau$, whereas shot noise models are generally restricted by the definition of special shot forms. Shot noise lightcurves therefore might be regarded as a subclass of LSSM lightcurves in the sense that a superposition of exponentially decaying shots can be interpreted as \\emph{one} possible realization of an intrinsic AR[1] process. The inspection of the measured lightcurves of Cyg~X-1 implies that the source of the derived AR[1] process is indeed the stochastic superposition of individual shot events, corresponding to the basic idea of shot noise. We note that on larger timescales this kind of variability is also present in the X-ray emission of active galactic nuclei (AGN) (e.g.~\\cite{mchardy:89,mushotzky:93,koenig:97a,koenig:97c}). The physical mechanism responsible for such a temporal behavior, however, is not yet understood. \\subsection{The LSSM results in the light of time-dependent Comptonization models}\\label{kapitel5.2} Recently, the discussion concerning accretion physics has begun to concentrate on the consideration of timing and spectral properties of the X-ray emission as two aspects of the same model (\\cite{kazanas:97a,koenig:97c,wilms:97c}; and references therein). The spectrum of both, AGN and the hard state of galactic black hole candidates, is usually explained by inverse Comptonization, where soft X-ray photons, provided by a cold accretion disk, are upscattered by inverse Compton collisions in a hot plasma to produce the observed high energy power-law. In this context it can be assumed that the observed X-ray temporal behavior is the result of the processing of short shots of seed photons within the accretion disk corona (\\cite{payne:80}). The initial photon pulse is hardened and temporally broadened while diffusing through the corona. As harder photons in the emerging spectrum have, on average, undergone more scattering events, they have stayed in the corona for a longer time than softer photons and the emerging pulse is comparatively broader (observed in AGN by \\cite{koenig:97c}). In addition, variability structures in high energy lightcurves exhibit a characteristic time-lag with respect to those in low energy lightcurves (Fig.~\\ref{fig6}). Frequency-dependent time-lags have e.g.\\ been observed in {\\sl Ginga} and RXTE observations of Cyg~X-1 (\\cite{miyamoto:88,miyamoto:92,wilms:97c}). The question to what extent these lags are produced in a hot corona or wether they are intrinsic properties of the cold disk is still unresolved (\\cite{miyamoto:88,miller:95,nowak:96,nowak:98}). \\citey{nowak:98}, however, show that the frequency-dependent time-lags of Cyg~X-1 can be qualitatively reproduced using a simple propagation model (but they also point out that the longest observed time-lags ($\\approx$0.1\\,sec) can probably not be explained by Comptonization). Another challenge for combined spectral and temporal theories is given by modelling the coherence of variability structures in different energy bands which is observed to be near unity over a large frequency range in Cyg~X-1 (\\cite{vaughan:97,wilms:97c}). \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{h0725f7.eps}} \\caption{Shot profiles for the energy-ranges 4--5\\,keV (dash-dotted line) and 12--23\\,keV (dashed line). The profiles have been computed for an accretion disk corona (ADC) model with an accretion geometry as described by \\protect\\citey{dove:97b}. This geometry is similar to an advection dominated flow (ADAF, \\protect\\cite{abramo:95,narayan:96}). The ADC model parameters were chosen to be appropriate for Cyg~X-1 (disk temperature: $kT_{\\rm d}$=200\\,eV, coronal temperature: $kT_{\\rm c}$=66\\,keV, optical depth of the corona: $\\tau_{\\rm c}$=2.1; \\protect\\cite{dove:97b}).}\\label{fig6} \\end{figure} Additional support for the time-dependent Comptonization model comes from the joint spectral and temporal analysis of AGN, which have similar properties to those of Cyg~X-1. For a sample of AGN observations we were able to show that a linear correlation between the photon index $\\Gamma$ and the LSSM[1] timescale $\\tau$ exists, in the sense that harder spectra have a longer variability timescale (\\cite{koenig:97c}). By comparing the observed $\\Gamma$-$\\tau$ relation with Monte Carlo simulations of time-dependent Comptonization models, it is possible to scale the model geometry (\\cite{koenig:97b}). For the hard-state galactic black hole candidates a similar study cannot be performed since not enough objects are known. What can be done, however, is to check whether our best-fit LSSM models for Cyg~X-1 could, at least roughly, be explained with a time-dependent Comptonization model. We have computed shot profiles for several possible accretion geometries using a linear Monte Carlo code. The detailed results of these simulations will be discussed elsewhere. Since shot noise can be regarded as a special representation of an AR[1] process (see Sect.~\\ref{kapitel5.1}), we can use the computed shots to generate lightcurves based on these profiles, including an additional white noise component. In analogy to our treatment of the distribution of $\\tau$ seen in the observations, we derive a most probable value for $\\tau$ from the simulated lightcurve samples, $\\tau_{\\rm th}$. We find $\\tau_{\\rm th}=(3\\ldots 6) R_{\\rm c}/c$, with $R_{\\rm c}$ being the radius of the corona. The profile of a Compton-shot only depends on the relative size of the system, parameterized by the light travel time through one radius of the spherical corona, $R_c/c$. By identifying $\\tau_{\\rm th}$ with the measured value of $\\tau=$0.19\\,s, it is therefore possible to express the measured $\\tau$ in terms of the coronal radius. Our simulations give a first estimate of 320--640 Schwarzschild radii for the coronal radius (assuming a mass of 10$M_\\odot$ for the black hole). Other authors estimate the size of the Comptonization cloud in the hard state to be $\\approx$100 Schwarzschild radii (\\cite{esin:97}, ADAFs) or $\\approx$23 Schwarzschild radii (\\cite{nowak:98}, from minimal time-lags). These values are only in rough agreement and a better understanding of the relation between spectral and temporal properties is needed to arrive at a consistent picture. We plan a more detailed study of the Compton-shot profiles as well as LSSM analyses of RXTE data to further constrain the accretion geometry of Cyg~X-1." + }, + "9802/astro-ph9802094_arXiv.txt": { + "abstract": "The evolution of the comoving cosmic merger-rate density of neutron star binaries $n_c(z)$ is calculated using a distribution of their merging times provided by population-synthesis computations of binary stars. We adopt an exponential law for the star formation rate with various timescales for different morphological types of galaxies. For elliptical galaxies also an initial burst of star formation, lasting one Gyr, is considered. The resulting $n_c(z)$ of most models agree with the form $n_c(z) \\propto (1+z)^{1.5-2}$ for $ z \\aplt 2$, which has been proposed for the source population of $\\gamma$-ray bursts. Assuming a standard candle luminosity, the computed brightness distribution is consistent with the BATSE results if bursts at the peak flux threshold, $P$~=~0.4 photons~cm$^{-2}$~s$^{-1}$, are located at a limiting redshift of 1.9 to 3.3. Progenitors of the systems producing $\\gamma$-ray bursts at small redshift (bright) are likely to host in spiral galaxies and star forming regions whereas these at high redshift (dim) reside mainly in elliptical galaxies. The location of a burst may be up to $\\sim 1$\\,Mpc away from the host galaxy. ", + "introduction": "The recent discoveries of optical counterparts of the $\\gamma$-ray bursts GRB~970228 (Groot \\etal\\ 1997) and GRB~970508 (Bond 1997) and the measured emission-line redshift of $z=0.853$ (Metzger et al.\\ 1997) for the latter provide evidence for their cosmological origin. A possible source for $\\gamma$-ray bursts (GRBs) may be the merging of two neutron stars in a close binary (Blinnikov \\etal\\ 1984). For cosmologies with no vacuum energy, the brightness distribution of burst intensities expected for a uniform source population is consistent with the BATSE distribution if the limiting redshift is about unity (e.g. Dermer 1992; Mao \\& Paczy\\'nski 1992). However, cosmological time-dilation effects in the BATSE sample indicate that the dimmest sources should be located at $z \\approx 2$ (Norris et al. 1995). If this is the actual limiting redshift, a source population with a comoving rate density of the form $n_c(z) \\propto (1+z)^{\\beta}$ is compatible with the BATSE distribution for $1.5 \\aplt \\beta \\aplt 2$ (Horack et al. 1995). Therefore, it is worthwhile to examine whether the comoving merger rate density of neutron-star binaries evolves in a similar fashion. By means of population synthesis computations for binary stars the merger rate of neutron-star binaries [hereafter \\nsns] can be computed (see Portegies Zwart \\& Yungelson 1998 and references therein). Using this approach Lipunov \\etal\\ (1995) computed the evolution of \\nsns\\ mergers as a function of redshift and $\\log N - \\log S$\\ distributions of GRBs for a cosmic population which contains galaxies with a constant star-formation rate and also galaxies with initial bursts of star formation in different proportions. Totani (1997) computed the evolution of the GRBs rate density from \\nsns\\ mergers in a model based on the observationally determined history of cosmic star formation and in a model derived from detailed galaxy evolution. Totani assumed a distribution of merging times $f(t_c)\\propto t_c^{-1}$ for \\nsns\\ systems\\footnote{Sahu \\etal\\ (1997) made similar computations for ``cosmic'' star formation history and a fixed merger time of $3 \\times 10^7$\\ yr.}. We go a step further by adopting for all galaxies exponentially decreasing star formation rates (SFR) with different timescales depending on galaxy morphology (\\eg\\ Sandage 1986). We apply the distribution of $t_c$ from model B of Portegies Zwart \\& Yungelson (1998, henceforth PZY98) which provides the best fit to the expected birthrate and orbital parameters of the Galactic population of high-mass binary pulsars. ", + "conclusions": "The computed relative merger rate as a function of redshift is in agreement with the detected rate of $\\gamma$-ray bursts up-to the limiting redshift of BATSE. For the models where elliptical galaxies experience an initial burst of star formation the occurrence rate of \\nsns\\ coalescence increases suddenly by more than a factor two at a redshift of 2.5, which is beyond our current detection limit. The synthetic $\\log N - \\log P$\\ distribution is compatible with the observations down to the completeness limit of $\\gamma$-ray catalogues. The same is true for models based on observationally inferred cosmic star formation history (see \\eg\\ Sahu \\etal\\ 1997) and for models based on galactic evolution (Totani 1997). They, however, predict a different behavior if star bursts occur at high $z$. The absolute value of the \\nsns\\ merger rate is found to be $~\\sim$~100 times larger than the GRBs frequency. Escape from this conundrum is obtained if the opening angle of the observed phenomenon is a few degrees, which is consistent with fire-ball models where leptons are converted into bulk barionic motion (M\\'esz\\'aros \\& Rees 1992). If $\\gamma$-ray bursts indeed originate from \\nsns\\ mergers and our model for star formation is correct, interesting implications follow. The progenitors of bright GRBs or bursts at low redshift ($z\\la 1.5$) most likely belonged to early-type spiral galaxies whereas the progenitors of the dimmest bursts and those at high redshift ($z\\ga 2$) were located in elliptical galaxies. The majority of the bursts of gravitational waves in this model are expected to originate from early-type spiral galaxies. If most stars in the Universe formed in dwarf star-burst galaxies, a substantial fraction of the parental population may originate from them. Suggested by current observational data the decline in the star formation rate beyond $z \\approx 1.2$~ (\\eg\\ Connolly \\etal\\ 1997), if real, will show-up as a turn over in the $\\log N - \\log P$\\ distribution for GRBs. As noticed by Tutukov \\& Yungelson (1994) and confirmed by PZY98 (their Figs. 6 and 8), space velocities of \\nsns\\ binaries may well exceed escape velocities of dwarf as well as giant galaxies and they may travel up to $\\sim 1$\\,Mpc before coalescence. Thus, a significant fraction of the sites of GRBs may not be directly associated with star forming regions." + }, + "9802/astro-ph9802027_arXiv.txt": { + "abstract": "A recalibration of the luminosity--linewidth technique is discussed which introduces (i) new cluster calibration data, (ii) new corrections for reddening as a function of inclination, and (iii) a new zero-point calibration using 13 galaxies with distances determined via the cepheid period--luminosity method. It is found that H$_0=82\\pm16$~km~s$^{-1}$~Mpc$^{-1}$ (95\\% confidence). ", + "introduction": "Good relative distances to galaxies can be found from the correlations between their global luminosities and rotation velocities (Tully \\& Fisher 1977). It is true that today there are several methods with smaller apparent scatter, such as exploitation of the planetary nebula luminosity function (Jacoby, Ciardullo, \\& Ford 1990), surface brightness fluctuations (Tonry, Ajhar \\& Luppino 1990), and type Ia supernova peak brightnesses (Reiss, Press, \\& Kirshner 1996). However the luminosity--rotation rate technique is one that can be applied to a large fraction of disk galaxies, hence thousands of cases all over the sky and out to substantial redshifts. It remains the most important distance tool for studies of {\\it deviations} from the universal expansion. It can give an estimate of the Hubble Constant with good $\\sqrt N$ statistics if given a zero-point calibration. To make the calibration, we look at galaxies that obey the correlations and which have independently known distances. There has been recent progress in the determination of distances to potential calibrators which makes it a good time to redefine the properties of the luminosity--linewidth correlations. The Hubble Space Telescope (HST) has been used to discover cepheid variable stars in 8 suitable galaxies and to determine distances based on the period--luminosity relations for those stars. Today there are 13 appropriate galaxies with distances determined with this method. In addition to the new zero-point calibration there are two other significant improvements to the methodology. First, more high quality data and more complete samples are available. Second, multicolor information extending to the infrared now provides better information about obscuration to luminosities as a function of inclination. ", + "conclusions": "" + }, + "9802/astro-ph9802357_arXiv.txt": { + "abstract": "{It is shown that in a number of scalar potentials with an unstable local maximum at the origin chaotic inflation is followed by new inflation if model parameters are appropriately chosen. In this model density fluctuation can have a large-amplitude peak on the comoving Hubble scale at the onset of the slow-roll new inflation and can result in formation of appreciable amount of primordial black holes on astrophysically interesting mass scales.\\\\ \\pacs{PACS Numbers: 98.80.Cq, 04.70.Bw. } \\newcommand{\\dw}{{\\rm DW}} \\newcommand{\\cw}{{\\rm CW}} \\newcommand{\\ml}{{\\rm ML}} \\newcommand{\\lt}{\\tilde{\\lambda}} \\newcommand{\\lh}{\\hat{\\lambda}} \\newcommand{\\phidot}{\\dot{\\phi}} \\newcommand{\\phicl}{\\phi_{cl}} \\newcommand{\\adot}{\\dot{a}} \\newcommand{\\phat}{\\hat{\\phi}} \\newcommand{\\ahat}{\\hat{a}} \\newcommand{\\hhat}{\\hat{h}} \\newcommand{\\phihat}{\\hat{\\phi}} \\newcommand{\\Nhat}{\\hat{N}} \\newcommand{\\hth}{h_{th}} \\newcommand{\\hbh}{h_{bh}} \\newcommand{\\gsim}{~\\mbox{\\raisebox{-1.0ex}{$\\stackrel{\\textstyle >} {\\textstyle \\sim}$ }}} \\newcommand{\\lsim}{~\\mbox{\\raisebox{-1.0ex}{$\\stackrel{\\textstyle <} {\\textstyle \\sim}$ }}} \\newcommand{\\bfx}{{\\bf x}} \\newcommand{\\bfy}{{\\bf y}} \\newcommand{\\bfr}{{\\bf r}} \\newcommand{\\bfk}{{\\bf k}} \\newcommand{\\bkp}{{\\bf k'}} \\newcommand{\\order}{{\\cal O}} \\newcommand{\\beq}{\\begin{equation}} \\newcommand{\\eeq}{\\end{equation}} \\newcommand{\\beqa}{\\begin{eqnarray}} \\newcommand{\\eeqa}{\\end{eqnarray}} \\newcommand{\\mpl}{M_{Pl}} \\newcommand{\\lmk}{\\left(} \\newcommand{\\rmk}{\\right)} \\newcommand{\\lkk}{\\left[} \\newcommand{\\rkk}{\\right]} \\newcommand{\\lnk}{\\left\\{} \\newcommand{\\rnk}{\\right\\}} \\newcommand{\\call}{{\\cal L}} \\newcommand{\\calr}{{\\cal R}} \\newcommand{\\half}{\\frac{1}{2}} \\newcommand{\\kc}{\\kappa\\chi} \\newcommand{\\bkc}{\\beta\\kappa\\chi} \\newcommand{\\gkc}{\\gamma\\kappa\\chi} \\newcommand{\\gbkc}{(\\gamma-\\beta)\\kappa\\chi} \\newcommand{\\dchi}{\\delta\\chi} \\newcommand{\\dphi}{\\delta\\phi} \\newcommand{\\dOmega}{\\delta\\Omega} \\newcommand{\\Phibd}{\\Phi_{\\rm BD}} \\newcommand{\\echi}{\\epsilon_\\chi} \\newcommand{\\ephi}{\\epsilon_\\phi} \\newcommand{\\Phihat}{\\hat{\\Phi}} \\newcommand{\\Psihat}{\\hat{\\Psi}} \\newcommand{\\that}{\\hat{t}} \\newcommand{\\Hhat}{\\hat{H}} \\newcommand{\\zk}{z_k} \\newcommand{\\msolar}{M_\\odot} ", + "introduction": "If overdensity of order of unity exists in the hot early universe, a black hole can be formed when the perturbed region enters the Hubble radius. While the properties of the primordial black holes (hereafter PBHs) thus produced were a subject of extensive study decades ago \\cite{Zel,Haw}, there were no observational evidence of their existence and only observational constraints were obtained against their mass spectrum \\cite{Carr,Nov}. Recently, however, possibilities of their existence have been raised from a number of astrophysical and cosmological considerations and there is an increasing interest in it. For example, they may be the origin of the massive compact halo objects (MACHOs) which are dark compact objects with typical mass $\\sim 1\\msolar$ and make up about $\\order (10^{-2})$ of the critical density \\cite{macho}. While the primary MACHO candidates are substellar baryonic objects such as brown dwarfs, it is difficult to reconcile such a large amount of these objects with the observed mass function of low mass stars \\cite{RF} and with the infrared observation of dwarf component \\cite{BU}, unless the mass function is extrapolated to the lower masses in an extremely peculiar manner of Population III stars are produced abundantly at the relevant mass scale \\cite{pop3}. Hence we should consider PBHs seriously as the second option. Furthermore this possibility can be experimentally tested by observing gravitational radiation from coalescing black hole MACHO binaries by laser interferometers \\cite{Naka}. Another interesting possibility is PBHs with mass $M\\sim 10^{15}$g which are just evaporating now \\cite{evaporate}. It has been argued that high energy phenomenon associated with evaporation can explain origin of a class of gamma-ray burst. For this purpose their abundance should be around $\\Omega=10^{-8}$ in unit of the critical density \\cite{Cline}. One may also consider formation of much heavier black holes with mass $M\\sim 10^8\\msolar$. Such black holes are expected to exist in the center of AGNs and quasars \\cite{Turner} and act as an engine of their activity. Although it is generally believed that these black holes are formed after recombination with a specially arranged spectrum of density fluctuations \\cite{Loeb}, it might be interesting to consider the possibility that they were of primordial origin, which might lead to a new scenario of galaxy formation. In any case, in order to produce PBHs on some specific scale, we must prepare density perturbation whose amplitude has a high peak of $\\order (10^{-2})$ on the corresponding scale. It is difficult, however, to realize such a spectral shape in inflationary cosmology \\cite{oriinf,newinf,chaoinf,inf}, which is not only indispensable to solve the horizon and the flatness problems but also the most sensible way to generate density perturbations \\cite{pert}, because usual models predict a scale-invariant spectrum. Given the smallness of the observed anisotropy of the cosmic microwave background radiation (CMB) \\cite{COBE}, the amplitude of primordial density fluctuation turns out to be $\\order (10^{-5})$ on any observable scales in most common models. It is possible, nevertheless, to realize a non scale-invariant spectrum by choosing somewhat contrived forms of the inflaton's potential. While examples of the various perturbation spectrum realized in different potentials is found in \\cite{Hodges}, we must admit that few of the models giving non scale-invariant spectrum with a single scalar field has a motivation in sensible particle physics. In particular, the model proposed by Ivanov et al.\\ \\cite{INN} in order to produce significant amount of PBHs employs a scalar potential with two breaks and a plateau in between. Several other toy potentials have also been studied in \\cite{BP}. One can construct more natural inflation models for PBH formation if one allows multiple scalar fields. Some of them make use of primordially isocurvature fluctuations \\cite{JY}, others include multiple stages of inflation with each regime governed by different field \\cite{ST,Ra,GL,Kawa}. We must assign appropriate form of coupling of the scalar fields in each model, which may not always be easy. \\newpage In the present paper, we propose a new scenario of multiple inflation which, unlike previous double inflation models \\cite{ST,Ra,GL,Kawa,di}, contains only one source of inflation. In this model we employ in the Einstein gravity an inflaton scalar field, $\\phi$, with a simple potential, $V[\\phi]$, which has an unstable local maximum at the origin, such as a quartic double-well potential. This is the same setup as the new inflation scenario \\cite{newinf}, but chaotic inflation is also possible if the $\\phi$ has a sufficiently large amplitude initially. In fact Brandenberger and Kung \\cite{BK} studied the initial distribution of a scalar field with such a potential and concluded that chaotic inflation is much more likely than new inflation. Thus we also start with chaotic inflation, but show that new inflation is also possible when $\\phi$ evolves towards the origin after chaotic inflation if the parameters of the potential is appropriately chosen so that the scalar field has the right amount of kinetic energy after chaotic inflation in order to climb up the potential hill near to the origin and start slow rollover there. Hence in this model the initial condition for new inflation is realized not due to the high-temperature symmetry restoration nor for a topological reason \\cite{topological}, but by dynamical evolution of the field which has already become sufficiently homogeneous because of the first stage of chaotic inflation. We shall refer this succession of inflation simply to chaotic new inflation. With an appropriate shape of the potential, density fluctuations generated during new inflation can have larger amplitude than those during chaotic inflation. Furthermore, since the power spectrum of fluctuation generated during new inflation can be tilted, it can have a peak on the comoving Hubble scale when the inflaton enters the slow-rollover phase during new inflation. If the peak amplitude is sufficiently large, it results in formation of PBHs on the horizon mass scale when the corresponding comoving scale reenters the Hubble radius during radiation domination. We shall show below that such a scenario is indeed possible with a simple shape of the inflaton's potential. The rest of the paper is organized as follows. In \\S II we summarize basic features of chaotic inflation and new inflation scenarios and in \\S III we calculate formation probability of PBHs in our model. Then we report the results of numerical analysis with various potentials in \\S IV and \\S V. Finally \\S VI is devoted to conclusion. ", + "conclusions": "In the present paper we proposed a new scenario of double inflation which contains only one inflation-driving scalar field, that is, we pointed out that a scalar potential which has an unstable local maximum at the origin can not only realize new and chaotic inflation separately but also accommodate both sequentially if its model parameters are appropriately chosen with natural initial conditions as employed in the chaotic inflation scenario. We have further shown that the spectrum of density fluctuation in this model can have a large-amplitude peak on the comoving Hubble scale at the onset of the slow-roll regime of new inflation and that this can be applied to formation of PBHs on a specific mass scale. This feature of the spectrum is realized naturally compared with other models with a single scalar field \\cite{INN,BP} because the scalar field is not slowly rolling at the onset of new inflation. On the other hand, we must specify the values of model parameters with many digits in order to produce an appropriate amount of PBHs on a desired scale. This feature, however, is more or less common to all the other models attempting to account for formation of PBHs in inflationary cosmology (see {\\it e.g.} \\cite{BP}), because both the peak amplitude of the fluctuations and its location must be specified with high accuracy due to the exponential dependence of the black hole abundance and its mass on the model parameters. On the other hand, one could in principle claim that observation of PBHs can serve as a strong tool to determine the parameters in the inflaton's dynamics. Unfortunately, however, our ignorance of the detailed condition for PBH formation, such as the precise threshold amplitude of fluctuation as a functional of the shape of perturbed region, makes it impossible to link the mass spectrum of PBHs with the shape of the inflaton's potential precisely. In the present paper we have calculated the values of the model parameters rather precisely under the universal assumption of $h_{bh}=0.75$. In fact, however, the values of the parameters would totally change had we chosen a different threshold. Hence the precise numbers we have quoted do not have much significance, but the number of digits simply indicates the sensitivity of the mass spectrum to the model parameters." + }, + "9802/astro-ph9802161_arXiv.txt": { + "abstract": "We propose a new method to constrain the actual state of the interstellar cloud that surrounds the solar system. Using Voyager UVS Lyman-$\\alpha$ sky maps and the powerful principle of invariance, we derive the H distribution all along the spacecraft path. Provided current models of the heliopause interface between the solar and the interstellar winds, we extrapolate this distribution to farther distances from the Sun and infer in a self consistent way key parameters of the local cloud. Our findings are a high interstellar hydrogen density of $\\sim0.24\\,{\\rm cm^{-3}}$ and a weak ionization $\\frac{{\\rm n(H^+)}}{{\\rm n(H^+)+n(H)}}\\simeq\\,14\\%$. ", + "introduction": "It is widely believed that the local interstellar medium represents a point sample of the more extended interstellar medium (ISM), and that its studies through its interplanetary signature is of particular interest to understand the origin and the evolution of the ISM. Because of the relative motion of the Sun with respect to the local interstellar cloud (LIC), the neutral component of the interstellar wind penetrates deeply inside the solar system, where it becomes accessible to in situ detections. After its discovery in the early 70's in the Earth neighborhood, interpretation of backscattered H Lyman-$\\alpha$ and HeI $58.4\\ {\\rm nm}$ solar radiations has been extensively employed to derive H and He abundances in the local cloud. As classical interpretations of UV sky maps usually require the modeling of the H distribution in the heliosphere with the problem of filtration at the heliopause and the calculation of the radiation field including multi-scattering effects, the derivation of neutral abundances proved to be difficult, and the values found were consequently poorly constrained.\\\\ The aim of this paper is to present a new technique for deriving the intrinsic properties of the LIC by interpreting differently the UV backscattered radiation measurements. This method consists in constraining the H distribution stage by stage, starting from the Earth neighborhood up to the far unperturbed LIC. In a first step, we determine the H distribution along the Voyager trajectory by applying the principle of invariance to UVS Lyman-$\\alpha$ sky maps recorded for several positions of the spacecraft (Puyoo et al 1997). Then we connect this distribution to the LIC one through a model of interaction between the solar and the interstellar flows. Lastly, assuming the LIC in a local steady state, we derive the intrinsic properties of the local cloud, i.e. the densities of both neutral and ionized Hydrogen and Helium in a self consistent way. ", + "conclusions": "After extraction of the hydrogen distribution from the Voyager Ly-$\\alpha$ sky maps inside the heliosphere (see Fig. 1) and assuming local statistical equilibrium, we derived that the LIC has a high H density ${\\rm n_{H\\infty}}\\simeq\\,0.24\\,{\\rm cm^{-3}}$ and is weakly ionized with an ionization fraction near the Sun of $\\sim\\,15\\%$ and $\\sim\\,6\\%$ for H and He respectively. The self-consistency of the method used here, makes the derived set of LIC parameters unique for the ionization sources considered here, which could however not explain the enhanced He ionization with respect to H found by Dupuis et al (1995).\\\\ In order to check the validity of the weak ionization inferred for H, it is interesting to compare our electron density to available measurements. We apply a simple model that calculates the photoionization of H and He from the Sun out to the edge of the local cloud to different lines of sight. Toward $\\epsilon$Cma or Capella, we obtain a mean electron density $<{\\rm n_e}>\\simeq\\,0.08\\,{\\rm cm^{-3}}$, which is in good agreement with recent HST measurements made by respectively Gry et al (1995) and Wood \\& Linsky (1995). As far as the ionization of H is the main pool for electron production, this result seems to confirm that this process is well described by our model. Nevertheless, an additional source that preferentially ionizes He atoms is required to explain EUVE results (Vallerga 1996)." + }, + "9802/astro-ph9802057_arXiv.txt": { + "abstract": "It is shown that for moderately hot polar caps (with effective temperature of $\\sim 10^6$ K), the efficiency of polar gap acceleration is lower compared to the case in which the polar caps are relatively cool and inverse Compton scattering plays no role in controlling the gap. For young pulsars with superstrong magnetic fields ($\\geq10^9\\ \\rm T$) and hot polar caps (with temperature of $\\ge 5\\times10^6\\,\\rm K$), because of the energy loss of electrons or positrons due to resonant inverse Compton scattering in the vicinity of polar caps, pair cascades occur at distances further away from the polar cap, and in this case, we have a relatively high acceleration efficiency with ions carrying most of the particle luminosity. ", + "introduction": "The mechanism for converting rotational energy of pulsars to electromagnetic radiation has long been the subject of active research in pulsar physics. The recent observation of seven pulsars by the instruments aboard CGRO provides even more challenge to the current pulsar theory. These pulsars radiate high-energy photons with apparently high efficiency (Ulmer 1994; Thompson et al. 1994). One possible way to convert the rotational energy to electromagnetic radiation is through particle acceleration by the rotation-induced electric field in the pulsar's magnetosphere. There are several sites for gap acceleration, which include the polar gap (Ruderman \\& Sutherland 1975; Arons \\& Scharlemann 1979; Michel 1974; Fawley, Arons \\& Scharlemann 1977), the slot gap (Arons 1983), and the outer gap (Cheng, Ho \\& Ruderman 1986). In the polar gap model, the acceleration region is located near the polar cap while in the outer gap model, the acceleration occurs near the light cylinder. In the slot gap model, the acceleration is at the boundary between the open and closed field lines. There is observational evidence that polar caps of young pulsars can be hot with effective temperature $T\\sim 10^6\\,\\rm K$ (\\\"Ogelman 1991; Greiveldinger et al. 1996). Although the exact value of $T$ is rather uncertain, and may depend on the model of the atmosphere above the polar cap (Romani 1987), the hot polar cap appears to be a plausible consequence of polar cap heating as the result of particle acceleration (Cheng \\& Ruderman 1977; Arons \\& Scharlemann 1979; Luo 1996). One of the important consequences of hot polar caps is that it provides an alternative mechanism for controlling the gap, i.e. mainly through pair cascades initiated by Compton scattered photons (Sturner 1995; Luo 1996). Inverse Compton scattering in the polar cap region is strongly modified by the magnetic field in that the scattering cross section is enhanced by the cyclotron resonance (Herold 1979), i.e. in the electron rest frame, soft photons have energies close to the cyclotron energy $\\varepsilon_B=B/B_c$ in $m_ec^2$ with $B_c\\approx4.4\\times10^9\\ \\rm T$ the critical field. This process is referred to as resonant inverse Compton scattering (RICS). Because of the resonance, the effective cross section is greatly enhanced in comparison with the ordinary Compton scattering. For gap acceleration, we can define its efficiency as the ratio of the total voltage across the gap to the maximum voltage across the polar cap (i.e. for an empty magnetosphere). The efficiency of the gap acceleration is controlled by the pair production and thus strongly depends on the effective temperature of the polar caps. In this paper, we consider the constraint imposed by RICS on the efficiency of polar gap acceleration. We assume that free emission of charges from polar caps is allowed, and consider both the polar region in which the outflowing charges are electrons (${\\bf B}\\cdot\\mbox{\\boldmath{$\\Omega$}}>0$) and that in which the outflowing charges are heavy ions or positrons (${\\bf B}\\cdot\\mbox{\\boldmath{$\\Omega$}}<0$). In the polar gap at which ${\\bf B}\\cdot\\mbox{\\boldmath{$\\Omega$}}<0$, since the energy loss of ions due to RICS is negligibly small, they can be accelerated by the full potential drop across the gap. We derive the condition for the ions to carry most of the particle luminosity. Accelerated ions may produce pairs in the anisotropic thermal photon field. We calculate numerically the distance from the polar cap at which the ions start to produce pairs and compare it to the gap length constrained by pair production by RICS. ", + "conclusions": "" + }, + "9802/astro-ph9802111_arXiv.txt": { + "abstract": "``Extreme Scattering Events'' (ESEs) are attributed to radio-wave refraction by a cloud of free-electrons crossing the line-of-sight. We present a new model in which these electrons form the photo-ionized `skin' of an underlying cool, self-gravitating cloud in the Galactic halo. In this way we avoid the severe over-pressure problem which afflicts other models. The UV flux in the Galactic halo naturally generates electron densities of the right order. We demonstrate, for the first time, a good reproduction of the prototypical ESE in the quasar 0954+658. The neutral clouds are a few AU in radius and have masses $\\lta10^{-3}\\, {\\rm M_\\odot}$. The observed rate of ESEs implies that a large fraction of the mass of the Galaxy is in this form. ", + "introduction": "ESEs were discovered ten years ago during radio flux monitoring of a sample of compact radio quasars (Fiedler et al 1987). The ESE phenomenon consists of dramatic flux changes occuring over several weeks to months. It is broadly agreed that ESEs are not intrinsic variations but, rather, apparent flux changes which are caused by refracting elements, a few AU in radius, crossing the line-of-sight. Both random (Fiedler et al 1987) and deterministic (Romani, Blandford \\& Cordes 1987) lens structures have been proposed; in all cases the refraction is attributed to free electrons. Two further points of consensus are that the blobs of free electrons must be Galactic, and that they represent a distinct component of the ISM (Narayan 1988; see also Rickett 1990). This same component is thought to be responsible for episodes of multiple imaging of radio pulsars (Cordes \\& Wolszczan 1986; Rickett 1990) --- a phenomenon which manifests itself as periodic fringes in the dynamic spectra. Despite the attention which the ESE phenomenon has attracted (Romani, Blandford \\& Cordes 1987; Romani 1988; Clegg, Chernoff \\& Cordes 1988; Clegg, Fey \\& Lazio 1998), the current state of understanding is unsatisfactory for two reasons. First the implied pressure of the electron cloud -- having an inferred density $n_e\\sim10^3\\;{\\rm cm^{-3}}$ and a temperature ${\\rm T\\sim10^4\\;K}$ -- exceeds that of the diffuse Interstellar Medium (ISM) by a factor of $10^3$. Such a cloud should dissipate on a time-scale comparable to the sound-crossing time, of order a year. This difficulty can be lessened by appealing to cloud geometries which are elongated along the line-of-sight, and by placing the clouds in regions of high pressure within the ISM --- specifically, old supernova shocks. In support of the latter idea it has been argued (Romani 1988) that the distribution on the sky of recorded ESEs reflects that of the strongest regions of non-thermal radio emission. However, with the small number of ESEs recorded to date, and the large area covered by the radio loops, this evidence is at best suggestive. The second main difficulty with existing models is that none has been able to reproduce the dual-frequency light-curve of the quasar 0954+658 --- the source for which the best ESE data exist. In this paper we present the essential elements of a new model for ESEs, wherein the free electrons comprise a `skin' of photo-ionized material around a cool, self-gravitating cloud. The pressure of the neutral material is balanced by the self-gravity of the cloud, while the ionized gas flows continuously from the surface; complete evaporation occurs over a time-scale of order the Hubble time. We show in \\S2 that this model reproduces the prototypical ESE light-curves. The event rate for ESEs then leads to the robust conclusion (\\S3) that these clouds contribute substantially to the Galactic dark matter. In \\S4 we note some key observational tests. ", + "conclusions": "We have presented a new model for the cause of Extreme Scattering Events, where radio-wave refraction by ionized material magnifies (``lenses'') a distant radio source. In our model the ionized material is generated by photo-evaporation of an underlying neutral, self-gravitating cloud, so there is no difficulty in understanding the requisite high electron pressures; indeed the ambient UV field in the Galactic halo naturally generates the necessary electron density. Moreover we find it straightforward to reproduce the dual-frequency light-curve of the prototypical ESE in the source 0954+658. These facts argue strongly for the soundness of the model. It follows immediately that a large fraction of the mass of the Galaxy is in the form of cool gas clouds of up to Jovian mass, and only a few AU in radius. We are unable to falsify this picture with existing data, but with new observations some straightforward tests can be made." + }, + "9802/astro-ph9802327_arXiv.txt": { + "abstract": "The longstanding faint blue galaxy problem is gradually subsiding as a result of technological advancement, most notably from high-resolution Hubble Space Telescope imaging. In particular two categorical facts have recently been established, these are: ~ \\noindent 1) The excess faint blue galaxies are of irregular morphologies, ~ and, ~ \\noindent 2) the majority of these irregulars occur at redshifts $1 < z < 2$. ~ \\noindent These conclusions are based on the powerful combination of morphological and photometric redshift data for all galaxies in the Hubble Deep Field to $I < 26$. Our interpretation is that the faint blue galaxy excess, which incidentally coincides with the peak in the observed mean galaxy star formation rate, represents the final formation epoch of the familiar spiral galaxy population. This conclusion is corroborated by the low abundance of normal spirals at $z > 2$. Taking these facts together we favour a scenario where the faint blue excess is primarily due to the formation epoch of spiral systems via merging at redshifts $1 < z < 2$. The final interpretation now awaits refinements in our understanding of the {\\it local} galaxy population. ", + "introduction": "The faint blue galaxy problem has been with us for several decades and is comprehensively reviewed in Koo \\& Kron (1992) and Ellis (1997). The problem is surmised as: {\\it an observed excess of faint galaxies over the zero or passive evolution model predictions}. This excess first arises at $b_{J} = 22$ mags and extends to the faintest magnitudes probed ($b_{J} = 28.5$ mags, c.f. Metcalfe {\\it et al.} 1995). The problem is compounded when one also considers the redshift distributions of galaxies at $b_{J}=22-24$ mags, e.g. Glazebrook {\\it et al.} (1995a), which in shape agree well with the zero-to-passive model predictions, but of course not in amplitude ({\\it i.e.} a reiteration of the original faint blue galaxy problem). Spectroscopic surveys to fainter magnitudes are currently limited by aperture and signal-to-noise considerations. The unfortunate situation then, is that the models require a continuous renormalisation to match the observations and such a solution often results in contrived and implausible physical implications. In this overview I summarise the recent substantial developments in the observational data from Hubble Space Telescope imaging and in particular the Hubble Deep Field (HDF), the current interpretation and finally the future observations required. The expectation is that through these refinements, as opposed to speculative retro-fitting, comes concordance. ", + "conclusions": "" + }, + "9802/hep-ph9802424_arXiv.txt": { + "abstract": "The $\\nu_\\mu$ and $\\nu_\\tau$ neutrinos (and their antiparticles) from a Galactic core-collapse supernova can be observed in a water-\\v{C}erenkov detector by the neutral-current excitation of $^{16}$O. The number of events expected is several times greater than from neutral-current scattering on electrons. The observation of this signal would be a strong test that these neutrinos are produced in core-collapse supernovae, and with the right characteristics. In this paper, this signal is used as the basis for a technique of neutrino mass determination from a future Galactic supernova. The masses of the $\\nu_\\mu$ and $\\nu_\\tau$ neutrinos can either be measured or limited by their delay relative to the $\\bar{\\nu}_e$ neutrinos. By comparing to the high-statistics $\\bar{\\nu}_e$ data instead of the theoretical expectation, much of the model dependence is canceled. Numerical results are presented for a future supernova at 10 kpc as seen in the SuperKamiokande detector. Under reasonable assumptions, and in the presence of the expected counting statistics, $\\nu_\\mu$ and $\\nu_\\tau$ masses down to about 50 eV can be simply and robustly determined. The signal used here is more sensitive to small neutrino masses than the signal based on neutrino-electron scattering. ", + "introduction": "When the core of a large star ($M \\ge 8 M_{\\odot}$) runs out of nuclear fuel, it collapses and forms a proto-neutron star with a central density well above the normal nuclear density (for a review of type-II supernova theory, see Ref.~\\cite{Bethe}). The total energy released in the collapse, i.e., the gravitational binding energy of the core ($E_B \\sim G_N M_ {\\odot}^2/R$ with $R \\sim$ 10 km), is about $3 \\times 10^{53}$ ergs; about 99\\% of that is carried away by neutrinos and antineutrinos, the particles with the longest mean free path. The proto-neutron star is dense enough that neutrinos diffuse outward over a timescale of several seconds, maintaining thermal equilibrium with the matter. When they are within about one mean free path of the edge, they escape freely, with a thermal spectrum characteristic of the surface of last scattering. The luminosities of the different neutrino flavors are approximately equal. Those flavors which interact the most with the matter will decouple at the largest radius and thus the lowest temperature. The $\\nu_\\mu$ and $\\nu_\\tau$ neutrinos and their antiparticles, which we collectively call $\\nu_x$ neutrinos, have only neutral-current interactions with the matter, and therefore leave with the highest temperature, about 8 MeV (or $\\langle E \\rangle \\simeq$ 25 MeV). The $\\bar{\\nu}_e$ and $\\nu_e$ neutrinos have also charged-current interactions, and so leave with lower temperatures, about 5 MeV ($\\langle E \\rangle \\simeq$ 16 MeV) and 3.5 MeV ($\\langle E \\rangle \\simeq$ 11 MeV), respectively. The $\\nu_e$ temperature is lower because the material is neutron-rich and thus the $\\nu_e$ interact more than the $\\bar{\\nu}_e$. The observation of supernova $\\nu_x$ neutrinos would allow the details of the picture above to be tested. For a detailed description of the supernova neutrino emission, including the justification of our choice of temperatures, see Refs.~\\cite{Woosley,Janka}. Even after many decades of experiments, it is still not known whether neutrinos have mass. Results from several experiments strongly suggest that neutrino flavor mixing occurs in solar, atmospheric, and accelerator neutrinos, and proof of mixing would be a proof of mass. The requirement that neutrinos do not overclose the universe gives a bound for the sum of masses of stable neutrinos (see \\cite{Raffelt} and references therein): \\begin{equation} \\sum_{i=1}^3 m_{\\nu_i} \\lesssim 100 {\\rm\\ eV}\\,. \\label{eq:cosmo} \\end{equation} However, direct kinematic tests of neutrino mass currently give limits for the masses compatible with the above cosmological bound only for the electron neutrino, $m_{\\bar{\\nu}_e} \\lesssim 5$ eV\\cite{Belesev}. For the $\\nu_\\mu$ and $\\nu_\\tau$ neutrinos the kinematic limits far exceed the cosmological bound: $m_{\\nu_\\mu} < 170$ keV\\cite{RPP}, and $m_{\\nu_\\tau} < 24$ MeV\\cite{RPP}. It is very unlikely that direct kinematic tests can improve these mass limits by the necessary orders of magnitude any time soon. As we will show in detail below, the most promising method for determining these masses is with supernova neutrinos. Even a tiny mass will make the velocity slightly less than for a massless neutrino, and over the large distance to a supernova will cause a measurable delay in the arrival time. A neutrino with a mass $m$ (in eV) and energy $E$ (in MeV) will experience an energy-dependent delay (in s) relative to a massless neutrino in traveling over a distance D (in 10 kpc) of \\begin{equation} \\Delta t(E) = 0.515 \\left(\\frac{m}{E}\\right)^2 D\\,, \\label{eq:delay} \\end{equation} where only the lowest order in the small mass has been kept. Since one expects one type-II supernova about every 30 years in our Galaxy\\cite{SNrate}, and since supernova neutrino detectors are currently operating, it is worthwhile to consider whether mass limits (or values) for $\\nu_x$ compatible with the cosmological bound, Eq.~(\\ref{eq:cosmo}), can be obtained. The problem of $\\nu_x$ mass determination with supernova neutrinos in existing (e.g., Refs.~\\cite{Wolfenstein,Seckel,Krauss,Burrows,Fiorentini}) and proposed detectors (e.g., Refs.~\\cite{Boron,SNBO,OMNIS}) has been considered before. The present work differs from the previous ones by the method with which the $\\nu_x$ are detected: inelastic scattering on $^{16}$O nuclei followed by proton or neutron emission, and subsequent gamma decay of excited $^{15}$N or $^{15}$O nuclei, as suggested in Ref.~\\cite{LVK}. We describe this signal and its time structure in Section II. In Section III we discuss the most relevant case of small masses. We find the smallest $\\nu_x$ mass that is recognizably different from zero in the presence of the expected finite counting statistics. In Section IV we show that the mass range is also limited from above. If the $\\nu_x$ mass is too large, the signal is broadened to such a degree that it disappears into the unavoidable background. We find the largest detectable $\\nu_x$ mass. Finally, in Section V we summarize our findings. ", + "conclusions": "We have presented a rather general method, including a thorough statistical analysis, of extracting information about the possible $\\nu_\\tau$ and $\\nu_\\mu$ masses from the future detection of a Galactic supernova neutrino burst by the SuperKamiokande detector. When such an event in fact occurs, the existing mass limits will be vastly improved and will approach, or cross over, the cosmological bound." + }, + "9802/astro-ph9802275_arXiv.txt": { + "abstract": "We use a continuous wavelet transform to analyze more than two decades of data for the BL~Lac object OJ~287 acquired as part of the UMRAO variability program. We find clear evidence for a persistent modulation of the total flux and polarization with period $\\sim 1.66$ years, and for another signal that dominates activity in the 1980s with period $\\sim 1.12$ years. The relationship between these two variations can be understood in terms of a `shock-in-jet' model, in which the longer time scale periodicity is associated with an otherwise quiescent jet, and the shorter time scale activity is associated with the passage of a shock; the different periodicities of these two components may reflect different internal conditions of the two flow domains, leading to different wave speeds, or different contractions of a single underlying periodicity, due to the different Doppler factors of the two flow components. We suggest that the modulation arises from a wave driven by some asymmetric disturbance close to the central engine. The periodic behavior in polarization exhibits excursions in $U$ which correspond to a direction $\\sim 45^{\\circ}$ from the VLBI jet axis. This behavior is not explained by the random walk in the $Q$-$U$ plane which is expected from models in which a pattern of randomly aligned magnetic field elements propagate across the visible portion of the flow, and suggests a small amplitude, cyclic variation in the flow direction in that part of the flow that dominates cm-wavelength emission. ", + "introduction": "It is generally accepted that centimeter-waveband emission from AGNs is associated with a jet of synchrotron plasma, the accretion structure and immediate environment of the central supermassive black hole contributing broadband emission from the infrared to $\\gamma$-rays (e.g., Bregman \\markcite{REF09} 1994). A number of processes may be responsible for temporal variations in the radio flux: the fueling (accretion) rate may change with time, leading to a long-term change in the jet power; the accretion disk may exhibit instability, influencing the extraction of energy in the form of a collimated outflow; the outflow itself may be Kelvin-Helmholtz unstable, leading to propagating internal structures (e.g. shocks); the outflow may interact with ambient inhomogeneities, which also may produce disturbances to the body of the flow (Wiita \\markcite{REF40} 1991; Birkinshaw \\markcite{REF06} 1991; Icke \\markcite{REF20} 1991). However, while there can be characteristic time scales associated with all such process, there is no a priori reason to believe that such variations will be {\\it periodic}. Nevertheless, a detailed understanding of blazars demands a detailed description of the temporal variations, which for completeness requires a search for periodicity. OJ~287 is a BL~Lac object with $z=0.306$ (Miller, French, \\& Hawley \\markcite{REF29} 1978; Sitko \\& Junkkarinen \\markcite{REF35} 1985), whose behavior differs in a number of intriguing ways from that of other blazars. Large amplitude, short time scale variations (Sillanpa\\\"a et al. \\markcite{REF34} 1994; Kidger \\& Gonz\\'alez-P\\'erez \\markcite{REF25} 1994; Valtaoja, Ter\\\"asranta, \\& Tornikoski \\markcite{REF38} 1994; Aller, Aller, \\& Hughes \\markcite{REF02} 1994) are seen in the light curves from the radio to optical wavebands. In the optical waveband, OJ~287 displays activity down to a time scale $\\sim 10$ minutes (Boltwood \\markcite{REF08} 1996). Numerous periodicities have been claimed, the strongest evidence being for a period $\\sim 12$ yr evident in optical data that span more than one century (Sillanpa\\\"a et al. \\markcite{REF33} 1988). This period has been interpreted by Valtonen and coworkers (e.g., Valtonen \\& Lehto \\markcite{REF39} 1997) as being due to the passage of the secondary black hole of a binary system through the accretion disk of the primary black hole. Interpreting VLB polarization observations in the context of a `shock-in-jet model', Cawthorne \\& Wardle \\markcite{REF11} (1988) suggests that the radio jet is aligned very close to the critical cone, $\\theta_c=\\sin^{-1}\\left(1/\\gamma\\right)$, in contrast with viewing angles of tens of degrees suggested by the modelling of events in some blazars (Hughes, Aller, \\& Aller \\markcite{REF17} 1989, \\markcite{REF18} 1991; Mutel et al. \\markcite{REF30} 1990; Carrara et al. \\markcite{REF10} 1993). This suggests that the flow should have a Doppler factor significantly in excess of unity, with a concomitant reduction in observed time scales -- making the source uniquely suited for an analysis of periodicity, because of the many cycles observed. Indeed, Katz \\markcite{REF24} (1997) has proposed a model based on a precessing disk, that would demand an extremely high Lorentz factor ($\\sim 50$), and thus a high Doppler factor if the flow is seen from close to, or within, the critical cone. Motivated by the preceding arguments, and by the intriguing claims of periodicity for the optical light curve of OJ~287, we performed a number of spectral analyses of centimeter waveband data for those sources well-enough observed to justify a structure function analysis (Hughes, Aller, \\& Aller \\markcite{REF19} 1992). Both Deeming \\markcite{REF12} (1975) and Scargle \\markcite{REF32} (1982) periodograms were constructed, together with smoothed power spectra using both standard tapering and smoothing techniques (Jenkins \\& Watts \\markcite{REF21} 1968), and a maximum entropy approach (Haykin \\& Kesler \\markcite{REF15} 1979; Ulrych \\& Ooe \\markcite{REF37} 1979). In general, each source exhibited a complex distribution of power with frequency, due in part to the fact that most activity is aperiodic, and in part to the observing window. In only one source -- OJ~287 -- was there some evidence for periodicity; in that case the Scargle false-alarm probability gave us some confidence in a variation with time scale $\\sim 1.6$ yr (Aller, Aller, Hughes, \\& Latimer \\markcite{REF03} 1992). Although a powerful technique in the right context, the fundamental limitation of Fourier methods is that they do not preserve the temporal locality of a signal. Power associated with the data's window will appear at low frequencies, possibly near to the frequency of a sought periodic signal, while outbursts with a character that is certainly not periodic will place power across the spectrum. Identifying periodicity that is well-hidden by other variations requires assessing the significance of one, among many, peaks in the power spectrum. A method of analysis that circumvents this problem is continuous wavelet analysis, which we discuss in detail in \\S~3. For a one-dimensional data set such as a time series, a wavelet analysis amounts to quantifying the behavior of the signal on different temporal scales, as a function of time. This is achieved by convolving the signal with a localized wave-packet, as the packet is translated along the series, for a number of `dilations' of the wave-packet, i.e., for progressively broader wave-packets, sensitive to progressively longer time scale behavior. Thus an initially narrow wave-form, sensitive to high frequency structure in the signal, is progressively stretched, or dilated, in time, making it sensitive to lower frequency structure. We are currently exploring the behavior exhibited by various sources in the UMRAO database using continuous wavelets, and here report results for OJ~287, having chosen this source as the first to study in light of its unique character. ", + "conclusions": "We have found that a continuous wavelet transform of the multifrequency, total flux and polarization data for the BL~Lac object OJ~287 clearly reveals a periodicity that is merely hinted at by a more conventional Fourier analysis. The modulation of the total flux persists for the duration of the time series available, but changes its period in the 1980s in a way that is consistent with the kinematics of a `shock-in-jet' model, wherein a portion of a previously quiescent flow is replaced by a propagating domain of enhanced emissivity. This may be a single event, or a series of closely-spaced individual events. In either case, the modulation appears to be an additional phenomenon, superposed on longer-term trends, and may be a consequence of a precessional-like motion of the flow in the vicinity of the $\\tau=1$ surface, as a consequence of propagating modes of the Kelvin-Helmholtz instability. Such behavior is evident also in the polarized flux, for which the most notable feature is a series of periodic excursions along a well-defined axis in the $Q$-$U$ plane which is at about $45^{\\circ}$ from the direction defined by the VLBI jet. This behavior differs in detail from what would be expected from a `random walk' model, and it also is suggestive of a periodic variation in the direction of the jet, and its mean magnetic field component, at that point along the flow from which most of the emission originates in the cm-waveband. Detailed modelling is called for to explore this behavior, but will require that allowance be made for projection and relativistic aberration effects. As the latter will be determined by the details of the flow's velocity field, it will be interesting to see if the observed behavior can be a generic feature of unstable flows, or requires finely-tuned geometry, magnetic field topology and velocity distribution. Were OJ~287 to be in an radio-active phase, intensive, space-based monitoring of the evolution of VLBI components recently ejected from the core would provide potentially exciting data, because such a component would define the detailed kinematics of the flow that dominates the total flux data. Unfortunately, as evidenced by UMRAO monitoring, the source is in a protracted quiescent phase, and further understanding of this object may have to await renewed activity. A possibility worth exploration is that OJ~287 is not unique amongst well-observed sources, but merely a source with particularly short time scale periodic components. It may well be that some long term variations seen in other sources are manifestations of a similar modulation, but expressed in those sources in a far less evident and accessible way: the UMRAO database is not temporally long enough to easily identify such behavior. We plan to explore other sources of the UMRAO database, in an attempt to identify similar activity through signatures such a phase coherence of outbursts over long times scales." + }, + "9802/astro-ph9802043_arXiv.txt": { + "abstract": "We have obtained the first broad-band X-ray spectra of the nearby compact elliptical galaxy M32 by using the {\\it ASCA} satellite. The extracted spectra and X-ray luminosity are consistent with the properties of the hard spectral component measured in giant elliptical galaxies believed to originate from X-ray binaries. Two {\\it ASCA} observations were performed two weeks apart; a 25\\% flux decrease and spectral softening occurred in the interval. We have also analyzed archival {\\it ROSAT} HRI data, and discovered that the X-ray emission is dominated by a single unresolved source offset from the nucleus of M32. We argue that this offset, combined with the extremely rapid large magnitude variations, and hard X-ray spectrum combine to weakly favor a (single) X-ray binary over an AGN origin for the X-rays from M32. The nuclear black hole in M32 must be fuel-starved and/or accreting from a radiatively inefficient advection-dominated disk: the product of the accretion rate and the radiative efficiency must be less than $\\sim 10^{-10}$ M$_{\\odot}$ yr$^{-1}$ if the X-ray source is indeed an X-ray binary. ", + "introduction": "At a distance of 700 kpc (\\cite{t90}), M32 (NGC 221) is by far the nearest elliptical galaxy to the Milky Way. Although extremely compact, with a half-light radius of $\\sim 110$ pc and correspondingly low velocity dispersion ($\\sigma\\sim 80$ km s$^{-1}$) and optical luminosity ($M_B=-15.7$), M32 is structurally ``normal'' (\\cite{k85}) and lies on the fundamental plane (\\cite{b92}). For these reasons it has played a foundational role in studies of stellar populations and kinematics in elliptical galaxies. In particular, M32 represents the earliest, and one of the strongest, dynamical cases for the presence of a massive ($\\sim 3\\ 10^6 M_{\\odot}$) black hole in a non-active elliptical galaxy (\\cite{b96}, \\cite{v97}, and references therein). M32 is also a source of X-rays. More than five magnitudes less optically luminous than any other elliptical galaxy detected in X-rays, the proximity of M32 presents a unique opportunity to push the investigation of the nature of X-rays from elliptical galaxies to intrinsically very faint systems and gain new insight into the ``hard component.'' It is now well-established (e.g., \\cite{f89}) that the X-ray emission from the brightest ellipticals is dominated by hot gas that originated as stellar mass loss and subsequently settled into hydrostatic equilibrium in the galactic potential. However, there is a theoretical expectation (e.g., \\cite{c91}) that less luminous galaxies have shallower gravitational potentials, so that below some transitional luminosity the gas becomes unbound and flows out in an unobservable galactic wind, thus revealing the presence of the integrated emission from X-ray binaries. And indeed, the hard X-ray flux from the ensemble of binaries -- circumstantial evidence of which was discovered from {\\it Einstein Observatory} observations (\\cite{c87}; \\cite{k92}) -- has been directly detected, first by BBXRT (\\cite{s93}), and then more universally by {\\it ASCA} (\\cite{a94}, \\cite{m94}, \\cite{m97}). While, as expected, the fraction of the X-ray flux contained within the hard component increases as the X-ray-to-optical flux ratio decreases, the soft component proves to be surprisingly persistent and generally dominates the spectrum at energies less than 2 keV where {\\it ASCA} and {\\it ROSAT} are most sensitive. In the 0.5-4.5 keV band, the luminosity of the hard component $L_{x,hard}\\sim 4.1\\ 10^{29} (L_B/L_{B,\\odot})$ erg s$^{-1}$, corresponding to a flux less than $5\\ 10^{-13}$ erg cm$^{-2}$ s$^{-1}$ or only $\\sim 70$ total (for four detectors, see next section) {\\it ASCA} source counts ks$^{-1}$ for a $10^{11} L_{B\\odot}$ galaxy at the distance of the Virgo Cluster (\\cite{m97}). Thus, the combination of the large distance of even the closest giant ellipticals, the intrinsic weakness of the hard component, and the obscuring effect of the ubiquitous hot gas component severely limits what constraints can be placed on the nature of the hard emission. Fits of thermal models to {\\it ASCA} spectra yield lower limits on the temperature of $\\sim 2$ keV for individual galaxy spectra and $\\sim 6.5$ keV for a composite spectrum; however, the data is equally well fit by a power-law with photon index $1.8\\pm 0.4$ (\\cite{m97}). It is not clear whether and, if so, to what extent an active nucleus (AGN) is contributing to the hard emission in some or all elliptical galaxies, as might be expected if they host ``dead'' quasars with some relic activity (\\cite{f95}). Because of the shallowness of its potential well, any gas bound to M32 will have $kT<0.1$ keV and an unimpeded view of the hard component is afforded. Despite this, previous X-ray studies have failed to resolve the binaries/AGN ambiguity in the origin of X-rays from M32, with the latest attempt being a detailed analysis of {\\it ROSAT} PSPC observations by \\cite{e96} (hereafter EWD96). In this paper we present the results of analyses of two recent {\\it ASCA} observations of M32, as well as of archival {\\it ROSAT} HRI data. Compared to the PSPC, {\\it ASCA} has superior spectral energy resolution and sensitivity at energies greater than 2 keV, enabling us to derive more accurate constraints and address the issue of spectral variability. The excellent spatial resolution of the HRI provides a considerable improvement in determination of the position and extent of the M32 X-ray emission. We have also examined archival {\\it Einstein Observatory} and {\\it EXOSAT} data, and additional insight is sought from inspecting the long term X-ray variability of M32. Although the evidence remains inconclusive, we argue that the X-ray emission is likely dominated by a single super-Eddington X-ray binary. If this is the case, the very low upper limit on the luminosity associated with accretion by the central black hole implies the presence of an advection-dominated disk and/or fuel-starved AGN with an extremely low value of the product of the accretion rate and radiative efficiency. ", + "conclusions": "While not definitively conclusive, the evidence points toward an X-ray binary rather than an AGN origin for the X-ray source in M32. The relatively short timescale and large magnitudes of the transient variations in flux (Figure 5) are more characteristic of X-ray binary than AGN behavior, and -- most significantly -- the X-ray emission region seems to be centered $\\sim 30$ pc from the center of M32. The X-ray source in M32 is unlikely to be a high mass X-ray Binary because of the old stellar population of M32. Among low mass X-ray Binaries (LMXBs), the X-ray temporal and spectral characteristics most resemble high luminosity non-transient neutron star LMXBs such as Sco X-1 or Cyg X-2 (\\cite{l95}). Another possibility is a binary containing a relatively massive ($>5$ M$_{\\odot}$) black hole, such as G2023+338. The extrapolation of the linear relationship between blue luminosity and the integrated X-ray luminosity from binaries leads to a predicted 0.5-4.5 keV luminosity of $\\sim 10^{38}$ erg s$^{-1}$ for M32 -- comparable to the average observed luminosity. However, the hard component is expected to be spatially distributed as the optical light if it consists of many unresolved, faint, discrete sources. Such an extended component would easily be detected by the {\\it ROSAT} HRI since it would produce $\\sim 200$ counts inside the half-light radius of $\\sim 40''$ -- about nine times the background. Any unresolved component cannot be much more luminous than $10^{37}$ erg s$^{-1}$. That the X-ray emission from M32 is more centrally concentrated than the light may simply be a consequence of the fact that, in elliptical galaxies, the integrated luminosity of discrete X-ray sources is dominated by the brightest ($\\sim 10^{38}$ erg s$^{-1}$) binary systems. If this is the case, it would not be surprising if, at the low optical luminosity of M32, a single binary dominates the X-ray emission and that it is located near the galactic nucleus where the overall stellar density is highly concentrated. The X-ray variability shown in Figure 5 implies that a single source dominates the emission, and not a compact population of discrete sources. If the X-ray emission were instead associated with an active nucleus (of an unusual kind), the ratio of integrated X-ray binary luminosity to optical luminosity would be lower than what has been measured in other observed spheroidal systems (\\cite{m97}). If the X-ray emission in M32 indeed originates in an X-ray binary, then the X-ray luminosity of any AGN is less than $5\\ 10^{36}$ erg s$^{-1}$ -- $\\sim 10^{-8}$ of the Eddington luminosity of the central massive black hole. (A conservative upper limit from the highest point in the long-term light curve -- see Figure 5 -- is only a factor of $\\sim 100$ greater.) This represents a uniquely low ratio of X-ray luminosity to black hole mass ($\\sim 3\\ 10^6 M_{\\odot}$; \\cite{v97}), compared with other objects in the sample of \\cite{h98}. This is consistent with the lack of AGN indicators at other wavelengths (EWD96), and implies that $\\epsilon_x\\dot M<10^{-10}$ M$_{\\odot}$ yr$^{-1}$, where $\\dot M$ is the accretion rate and $\\epsilon_x$ the efficiency of converting fuel into X-rays. Evidently, the black hole is fuel-starved and/or has a highly radiatively inefficient accretion disk. The integrated mass loss from stars in M32 is $\\sim 0.005$ M$_{\\odot}$ yr$^{-1}$. Since the central inflow rate would be of this order for a global cooling flow, this would seem to argue against a lack of sufficient available fuel for the central black hole. However, until there are direct observational limits on the central gas density in the relevant 50-100 eV temperature range, this possibility cannot be excluded. Radiatively inefficient, advection-dominated accretion has been suggested as an explanation for the relatively low X-ray luminosities in LINERs (e.g., NGC 4258, \\cite{l96}), and for the paucity of bright nuclear X-ray sources in elliptical galaxies that statistical arguments suggest should host $10^8$-$10^9$ M$_{\\odot}$ black holes (\\cite{f95}, \\cite{ma97}). The ratio of X-ray luminosity to black hole mass in M32 is greater than 1000 times lower than in NGC 4258 if, as we have argued, the AGN is not the primary X-ray source. The corresponding accretion rate for an advection dominated disk with the `standard' structure parameters of \\cite{ma97} is less than $2.5\\ 10^{-4}(\\eta_x/0.1)^{-1/2}$ in Eddington units, where $\\eta_x$ is the fraction of the disk luminosity emitted in the {\\it ROSAT} band. Since this corresponds to $\\dot M<1.6\\ 10^{-5} (\\eta_x/0.1)^{-1/2}$ M$_{\\odot}$ yr$^{-1}$, the M32 black hole cannot be fueled at the rate expected for a global cooling flow unless $\\eta_x<10^{-6}$. The low luminosity and accretion rate in M32 is more in line with the advection-dominated accretion disk model for the putative black hole in the center of our own galaxy (\\cite{n95}) than with NGC 4258 or other LINERS." + }, + "9802/gr-qc9802004_arXiv.txt": { + "abstract": "\\mbox{}\\\\ The small or zero cosmological constant, $\\Lambda$, probably results from a macroscopic cancellation mechanism of the zero-point energies. However, nearby horizon surfaces any macroscopic mechanism is expected to result in imperfect cancellations. A phenomenological description is given for the residual variable cosmological constant. In the static, spherically symmetric case it produces approximate black holes. The model describes the case of exponential decay by $\\Box\\ln\\Lambda=-3a$, were $a$ is a positive constant. ", + "introduction": " ", + "conclusions": "" + }, + "9802/astro-ph9802039_arXiv.txt": { + "abstract": "Results of numerical simulations of the impact of a common envelope on the matter flow pattern near the outflowing component in a semidetached binary system are presented. Three-dimensional modeling of the matter transfer gas dynamics in a low-mass X-ray binary X1822--371 enable investigation of the structure of flows in the vicinity of the inner Lagrange point $L_1$. Taking into account the common envelope of the system substantially changes the flow pattern near the Roche surface of the outflowing component. In a stationary regime, accretion of common envelope gas is observed over a significant fraction of the donor star's surface, which inhibits the flow of gas along the Roche surface to $L_1$. The change in the flow pattern is particularly significant near $L_1$, where the stream of common envelope gas strips matter off the stellar surface. This, in turn, significantly increases (by an order of magnitude) the gas flow from the donor surface in comparison with the estimates of standard models. ", + "introduction": "The observational manifestations of interaction in semidetached binary systems (cataclysmic binaries, low-mass X-ray binaries, and supersoft X-ray sources) are extremely interesting. Comparisons of computational results from numerical simulations of matter flows in these systems with observational data provide a basis for studying and understanding the physical processes occurring in these binaries. As a rule, numerical simulations of semidetached systems have been performed using the generally well-justified assumptions that the orbits of the system components are circular and the components' rotation is synchronous with their orbital motion. In this case, the standard scenario for mass transfer in the system is as follows. In the course of its evolution, the donor star fills its critical surface (which, under the outlined assumptions, coincides with the Roche surface in the restricted three-body problem), and mass begins to flow through the vicinity of the inner Lagrange point $L_1$, where the pressure gradient is not balanced by the gravitation force. The flow pattern in semidetached systems has been investigated using both analytical [1--3] and numerical [4--15] models. However, in all these studies, the impact of the common envelope of the system on the structure of gaseous flows was either not taken into account at all, or was taken into account not entirely correctly. The morphology of flows in binary systems allowing for the presence of a common envelope was first considered by us in [16] (henceforth, Paper I). There, as well as in this paper, the \"common envelope\" of the system refers to the gas filling the space between the two components of the system and not involved in the accretion process (i.e., not belonging to the accretion disc). Our results of three-dimensional numerical simulations of a low mass X-ray binary X1822--371 in Paper I showed the contribution of the common envelope to the formation of flow structures in the system to be significant. The presence of a common envelope also influences the flow pattern in the vicinity of the donor star, which, in turn, is reflected in the mass exchange rate in the system. This problem is extremely important, because though the observational manifestations of interaction primarily depend on the general flow pattern studied in Paper I, the evolution of the system is determined by the mass exchange parameters. Here, we present the results of an analysis based on the calculations described in Paper I, but with the main accent on a detailed consideration of the influence of the common envelope on the structure of gas flows near the Roche surface of the outflowing component. ", + "conclusions": "Analysis of the three-dimensional numerical simulation results presented here indicates that taking into account the impact of the common envelope in semidetached binary systems radically changes the mass exchange parameters obtained. For the low-mass X-ray binary X1822--371 that we have considered, the overall mass inflow to the system increased by a factor of about 50 compared to the estimates of standard models for the same gas parameters values at the surface of the donor star. In addition, the common envelope gas changed the flow pattern near the surface of the outflowing component, which ultimately affected the overall structure of the gas flows in the system, and, thus, the expected observational manifestations of these flows. Unfortunately, the quantitative estimates we have obtained are valid only for the specific object considered, and do not allow us to draw more general conclusions about changes in the mass exchange parameters expected for other semidetached systems. This is due to the fact that the mass flow increase depends not only on the parameters of the binary system, but also on the properties of the common envelope gas, which can only be determined using three-dimensional numerical simulations. Nevertheless, given the results presented here, it is clear that an adequate description of the mass flows in any semidetached system with a common envelope is possible only in models that take into consideration the influence of this envelope." + }, + "9802/astro-ph9802349_arXiv.txt": { + "abstract": "Recently, it has been claimed that the recurrent nova T Pyx exhibits oppositely directed jets of ejecta apparent in features seen in H$\\alpha$ emission. Here we demonstrate that these features are in fact emission in the [N{\\scriptsize II}] lines which lie either side of H$\\alpha$ and arise from the expanding shell associated with this object rather than from collimated jets. We estimate an expansion velocity along a line of sight through the centre of the shell of about 500 km~s$^{-1}$. ", + "introduction": "T Pyx is a recurrent nova with recorded outbursts in 1890, 1902, 1920, 1944 and 1966 (\\cite{w87}). It is notable for possessing a bright nebular shell extensively investigated by Shara et al.\\ (1989) and Shara et al.\\ (1997). Recently, Shahbaz et al.\\ (1997 -- hereafter S97) have presented optical spectroscopy of T Pyx in which they identify emission components to the red and blue of H$\\alpha$ (S$^+$ and S$^-$ respectively in their Figure 1) which they interpret as red- and blue-shifted H$\\alpha$ emission from oppositely directed jets. These features occur at 6593 and 6539~$\\rm{\\AA}$ respectively implying line-of-sight velocities of 1380 and $-1082$ km~s$^{-1}$. ", + "conclusions": "On the night of November 21-22 1997 we obtained several spectra of T Pyx with LRIS (the Low Resolution Imaging Spectrograph, Oke et al.\\ 1995) on the Keck II telescope on Hawaii. Two 400~second exposures were made at slit position angles of 30$^\\circ$ and 120$^\\circ$. The 600 grooves~mm$^{-1}$ grating was used with a slit width of 1.5~arcsec matching the seeing. The pixel scale was 0.2~arcsec and the spectral resolution 8.8~${\\rm \\AA}$. Figure 1 shows a greyscale representation of the 2-dimensional spectrum from position angle 30$^\\circ$ clearly revealing the presence of an expanding shell emitting in [N{\\scriptsize II}] and H$\\alpha$ with signal to noise ranging from 10 to 70. The spectrum from position angle 120$^\\circ$ is similar although because the shell is not spherically symmetric (see Shara et al.\\ 1997) there are detailed differences. In order to relate these data to the spectra presented by S97, Figure 1 also includes the spectrum obtained by summing along the slit. Although the spectral resolution is lower than that of S97 it is still obvious that their features S$^+$ and S$^-$ are in fact, respectively, the red-shifted component of [N{\\scriptsize II}]~6583 and the blue-shifted component of [N{\\scriptsize II}]~6548. Further to this, by extracting the stellar continuum from our 2D spectrum using Horne's optimal extraction method available in the software package Figaro, and then subtracting this from the total spectrum shown in Figure 1, we obtain a reasonable approximation to the shell-only summed spectrum -- see Figure 2. In this spectrum we have indicated the rest wavelengths of H$\\alpha$ and the two [N{\\scriptsize II}] lines together with the positions of the blue- and red-shifted components arising from the front and back of a shell expanding at a velocity of 530~km~s$^{-1}$. It is worth noting that the component at about 6573~${\\rm \\AA}$ (arising from a combination of blue-shifted [N{\\scriptsize II}]~6583 and red-shifted H$\\alpha$) is clearly visible in the spectrum in Figure 1 of S97, whilst the 6555~${\\rm \\AA}$ component (from red-shifted [N{\\scriptsize II}]~6548 and blue-shifted H$\\alpha$) is blended into the blue wing of H$\\alpha$ in their Figure 1. It is difficult to estimate the expansion velocity from these spectra. Apart from the problems of contamination between H$\\alpha$ and [N{\\scriptsize II}], the shell is clumpy and incomplete. The velocity of 530~km~s$^{-1}$ is derived from the wavelength 6595~${\\rm \\AA}$ of the red-shifted [N{\\scriptsize II}]~6583 line at the point where it crosses the stellar continuum. This corresponds to the expansion velocity along a line of sight through the centre of the shell and hence will be independent of slit position angle. We estimate an uncertainty of $\\pm 2$~${\\rm \\AA}$ on this wavelength, equivalent to $\\pm 90$~km~s$^{-1}$, as a result of the contamination by the stellar continuum and of the spectral resolution. A better estimate would require more kinematical data at higher resolution across the whole shell and a plausible model for its structure. Note that the summed spectrum shown in Fig.\\ 2 peaks shortward of 6595~${\\rm \\AA}$ because this particular feature is dominated by emission from a bright part of the shell 2-3~arcsec below the star (see Fig.\\ 1) which is at lower radial velocities. In conclusion we suggest there is little evidence to support the existence of collimated jets in T Pyx. The data are however consistent with the presence of a shell expanding at about 500~km~s$^{-1}$ and emitting more strongly in [N{\\scriptsize II}] than in H$\\alpha$ -- this is in broad agreement with the findings of Shara et al.\\ (1989). However, more detailed modelling of the structure of the shell and further higher spectral resolution observations are required in order to reconcile this line-of-sight expansion velocity with the upper limit on the velocity in the plane of the sky of 40~km~s$^{-1}$ derived from Hubble Space Telescope observations of the proper motion of knots in the nebular shell (\\cite{s97})." + }, + "9802/astro-ph9802312_arXiv.txt": { + "abstract": " ", + "introduction": "To date, calculations of neutrino opacities in dense matter have received relatively little attention~\\cite{Saw,IP,HW,RPL} compared to other physical inputs such as the equation of state (EOS). The neutrino cross sections and the EOS are intimately related. This relationship is most transparent in the long-wavelength or static limit, in which the response of a system to a weak external probe is completely determined by the ground state thermodynamics (EOS). Thus, in this limit, neutrino opacities consisitent with the EOS can be calculated~\\cite{Saw}. However, when the energy and momentum transfered by the neutrinos are large, full consistency is often difficult to acheive in practice. Despite this, many salient features associated with an underlying dense matter model may be incorporated in the calculation of the neutrino opacities. In \\S2, we describe how this is accomplished for a non-relativistic potential model. The effects of nucleon-nucleon correlations on the neutrino mean free paths are calculated using the random phase approximation (RPA) in \\S3, where we show that the magnitudes of these many-body effects on both the scattering and absorption reactions are large. The total scattering cross section in a multi-component system including the effects of correlation due to both strong and electromagnetic interactions are presented in \\S4 using a relativistic formalism. The implications of these results for neutrino transport in a protoneutron star are in \\S5. ", + "conclusions": "We have highlighted the influence of correlations and collective phenomena on the neutrino opacities in dense matter. Our findings here indicate that the neutrino cross sections are significantly reduced, and the average energy transfer in neutrino-nucleon interactions is increased due to the presence correlations in the medium. Several improvements are necessary before we can assess the influence of these results on the macrophyscial evolution of a protoneutron star. Among the most important of these are calculations that provide (1) the dynamic form factor, (2) the particle-hole and particle-particle interactions, (3) the renormalization of the axial charge, and, (4) the means to assess the role of multi-pair excitations, in charge neutral, beta-equilibrated dense matter at finite temperature. While investigations along these directions are in progress (see also \\cite{BS}), some general trends may be anticipated. The many-body effects studied here, including the improvements listed above, suggest considerable reductions in the opacities compared to the free gas estimates often employed in many applications. In the particular instance of the early evolution of a protoneutron star, the suggested modifications imply shorter time scales over which the deleptonization and cooling occur. Since mean free paths $\\lambda$ are much less than the stellar radius $R$, evolution is via diffusion. To order of magnitude, therefore, the timescale $\\tau\\propto1/\\lambda$. However, there are important feedbacks between these evolutionary consequences and the underlying EOS, in particular the specific heat of multicomponent matter, which have to be studied before firm conclusions may be drawn." + }, + "9802/gr-qc9802025_arXiv.txt": { + "abstract": "A Green's function method is developed for solving strongly-coupled gravity and matter in the semiclassical limit. In the strong-coupling limit, one assumes that Newton's constant approaches infinity, $G \\rightarrow \\infty$. As a result, one may neglect second order spatial gradients, and each spatial point evolves like an homogeneous universe. After constructing the Green's function solution to the Hamiltonian constraint, the momentum constraint is solved using functional methods in conjunction with the superposition principle for Hamilton-Jacobi theory. Exact and approximate solutions are given for a dust field or a scalar field interacting with gravity. ", + "introduction": "The action for Einstein gravity interacting with a dust field, $\\chi$, and a scalar field, $\\phi$, may be written as follows: \\numparts \\begin{equation} {\\cal I} = \\int d^4x \\, \\sqrt{ -g} \\; \\left \\{ {1 \\over 2 \\kappa } {}^{(4)}R - {1 \\over 2 \\kappa } g^{\\mu \\nu} \\phi_{,\\mu} \\phi_{,\\nu} - \\kappa V(\\phi) - {n \\over 2} ( g^{\\mu \\nu} \\chi_{,\\mu} \\chi_{,\\nu} + \\kappa^2 ) \\right \\} \\, , \\label{action} \\end{equation} where \\begin{equation} \\kappa \\equiv 8 \\pi G = 8 \\pi/ m_P^2 \\, \\end{equation} \\endnumparts is the gravitational coupling constant. $n \\equiv n(t,x)$ is a Lagrange multiplier which ensures that the square of the four-velocity, $U^{\\mu}$, is minus one: \\begin{equation} U^\\mu \\, U_\\mu = - 1 \\, , \\quad U^\\mu = - g^{\\mu \\nu} \\, \\chi_{,\\nu} / \\kappa\\, . \\end{equation} The above form can be obtained from the usual one by appropriate scalings of the matter fields and their coupling constants. Factors of $\\kappa$ appear in key places of eq.(\\ref{action}) in order to obtain the desired energy constraint: \\numparts \\begin{eqnarray} {\\cal H}(x) &&= \\kappa \\; \\gamma^{-1/2} \\, \\left ( 2 \\gamma_{ac} \\gamma_{bd} - \\gamma_{ab} \\gamma_{cd} \\right ) \\pi^{ab} \\pi^{cd} + \\kappa { \\gamma^{-1/2} \\over 2 } \\left ( \\pi^{\\phi} \\right )^2 + \\nonumber \\\\ && \\kappa \\; \\pi^\\chi \\sqrt{ 1 + \\chi^{|a} \\chi_{|a}/ \\kappa^2 } -{\\gamma^{1/2} \\over 2 \\kappa} \\; R + {\\gamma^{1/2} \\over 2 \\kappa} \\phi^{|a} \\phi_{|a} + \\kappa \\; \\gamma^{1/2} \\, V(\\phi) = 0\\; . \\label{ec} \\end{eqnarray} No factors of $\\kappa$ appear in the momentum constraint: \\begin{equation} {\\cal H}_{i}(x)=-2\\left(\\gamma_{ik} \\pi^{kj} \\right)_{,j} + \\pi^{kl} \\gamma_{kl,i} + \\pi^\\phi \\phi_{,i} + \\pi^\\chi \\chi_{,i} = 0 \\, . \\label{mc} \\end{equation} \\endnumparts In a classical Hamilton-Jacobi formulation of general relativity, one defines the generating functional, \\numparts \\begin{equation} {\\cal S} \\equiv {\\cal S}[\\gamma_{ab}(x), \\phi(x), \\chi(x)], \\end{equation} by assigning a real number to each field configuration $[\\phi(x),\\chi(x)]$ on a space-like hypersurface with 3-geometry given by $\\gamma_{ab}(x)$. In a semiclassical context, one allows for the possibility that ${\\cal S}$ may be complex. The Hamilton-Jacobi equations are obtained from the constraints (\\ref{ec}-b) by replacing the momenta with functional derivatives of ${\\cal S}$: \\begin{equation} \\pi^{\\chi}(x) = { \\delta {\\cal S} \\over \\delta \\chi(x) }, \\quad \\pi^{\\phi}(x) = { \\delta {\\cal S} \\over \\delta \\phi(x) }, \\quad \\pi^{ab}(x) = { \\delta {\\cal S} \\over \\delta \\gamma_{ab}(x) } \\, . \\end{equation} \\endnumparts In the limit of large gravitational coupling, $\\kappa \\rightarrow \\infty$, the generating functional ${\\cal S}^{(s)}$ for the strongly-coupled system evolves according to the following equations: \\numparts \\begin{eqnarray} {\\cal H}^{(s)}(x)/ \\kappa = && \\; \\gamma^{-1/2}\\, \\left ( 2 \\gamma_{ac} \\gamma_{bd} - \\gamma_{ab} \\gamma_{cd} \\right ) {\\delta {\\cal S}^{(s)} \\over \\delta \\gamma_{ab}} {\\delta {\\cal S}^{(s)} \\over \\delta \\gamma_{cd}} + \\; {\\delta {\\cal S}^{(s)} \\over \\delta \\chi} \\nonumber \\\\ && + {\\gamma^{-1/2} \\over 2 } \\left ( {\\delta {\\cal S}^{(s)} \\over \\delta \\phi} \\right )^2 + \\; \\gamma^{1/2} \\, V(\\phi) \\; = 0 \\; , \\label{ecs} \\end{eqnarray} \\begin{equation} {\\cal H}^{(s)}_{i}(x)=-2\\left(\\gamma_{ik} \\, {\\delta {\\cal S}^{(s)} \\over \\delta \\gamma_{kj}} \\right)_{,j} + {\\delta {\\cal S}^{(s)} \\over \\delta \\gamma_{kl}} \\gamma_{kl,i} + { \\delta {\\cal S}^{(s)} \\over \\delta \\phi } \\, \\phi_{,i} + { \\delta {\\cal S}^{(s)} \\over \\delta \\chi } \\, \\chi_{,i} = 0 \\, . \\label{mcs} \\end{equation} \\endnumparts The energy constraint is ultra-local in the sense that different spatial points are not coupled to each other. As a result, the Poisson bracket of ${\\cal H}^{(s)}(x)$ with ${\\cal H}^{(s)}(y)$ vanishes, and consistency of the Hamiltonian constraint at different spatial points is assured. Consistency or `integrability' of the Hamiltonian constraint for full general relativity is more complicated. It is related to the freedom in choosing an arbitrary time foliation: see Parry \\etal (1994) as well as Salopek (1995). However, the momentum constraint does couple the spatial points together because spatial derivatives appear. It is this cross-coupling of different spatial points that makes the strongly-coupled system non-trivial. (In order to simplify the notation, ${\\cal S}$ will be used to denote ${\\cal S}^{(s)}$ for the remainder of this paper.) \\subsection{Analogy from Elementary Quantum Mechanics} The general solution for the strongly-coupled system may appear quite complicated so it is instructive to consider a well-known example from elementary quantum mechanics which illustrates the essential features. \\subsection{Rotationally Symmetric Solutions to the Schrodinger Equation for a Free Particle} The two-dimensional Schrodinger equation for a free particle with zero angular momentum is given by: \\numparts \\begin{equation} i { \\partial \\psi \\over \\partial t} = - { 1 \\over 2 m} \\left ( {\\partial^2 \\psi \\over \\partial x^2} + {\\partial^2 \\psi \\over \\partial y^2} \\right ) \\, , \\label{schrod} \\end{equation} \\begin{equation} L_z \\; \\psi = i \\left ( y { \\partial \\over \\partial x} - x { \\partial \\over \\partial y} \\right ) \\psi = 0 \\, . \\label{symmetry} \\end{equation} \\endnumparts In solving these two equations, one would ordinarily invoke polar coordinates, $(r, \\theta)$, and then write the solution as $\\psi \\equiv \\psi(r)$. Instead, I will utilize a circuitous method which is of interest because the same technique may be generalized to the case of strongly-coupled gravity. One notes immediately that the first equation (\\ref{schrod}) can be solved generally without any reference to the second by using a superposition of plane waves: \\begin{equation} \\Psi = \\int d^2k \\; f(\\vec k) \\; e^{i ( \\vec k \\cdot \\vec x - \\omega t ) } \\, \\quad w = (k_1^2 + k_2^2)/ (2m), \\end{equation} where $f(\\vec k)$ is an arbitrary function of the wave-vector $\\vec k$. Please note that a plane wave $e^{i ( \\vec k \\cdot \\vec x - \\omega t ) }$ is {\\it not} a solution of the symmetry condition, eq.(\\ref{symmetry}). However, by suitably restricting the function $f$ so that it is function only of the magnitude of $k = \\sqrt{k_1^2 + k_2^2}$, \\begin{equation} f(\\vec k) \\equiv f(k) \\, \\end{equation} one can indeed satisfy the symmetry condition. In a semiclassical gravitational context, the energy constraint (\\ref{ecs}) will be analogous to the Schrodinger equation (\\ref{schrod}) and the momentum constraint (\\ref{mcs}) will be analogous to the symmetry condition (\\ref{symmetry}). In general relativity, it is not yet known how to explicitly solve the momentum constraint such that one obtains a set of reduced variables ({\\it physical variables}) analogous to solving the $L_z$ constraint and then deducing that $\\psi$ is solely a function of $r$. Instead, one solves the strongly-coupled system by first obtaining a general class of solutions for the energy constraint. By considering a suitable superposition over these solutions, one then constructs a solution to the momentum constraint. Shortly after Dirac (1958) gave general relativity its Hamiltonian form, Higgs (1958) pointed out that the momentum constraint of general relativity implied that the wavefunctional was invariant under spatial coordinate transformations. Although, the Hamilton-Jacobi equation was written explicitly by Peres (1962), only quite recently (Salopek and Stewart 1992, Parry \\etal 1994) has the spatial coordinate invariance been exploited to construct explicit solutions of the full HJ equation. They utilized a spatial gradient expansion. ", + "conclusions": "" + }, + "9802/astro-ph9802062_arXiv.txt": { + "abstract": "\\noindent The kaon energy in a nuclear medium and its dependence on kaon-nucleon and nucleon-nucleon correlations is discussed. The transition from the Lenz potential at low densities to the Hartree potential at high densities can be calculated analytically by making a Wigner-Seitz cell approximation and employing a square well potential. As the Hartree potential is less attractive than the Lenz one, kaon condensation inside cores of neutron stars appears to be less likely than previously estimated. ", + "introduction": "Kaon condensation in dense matter was suggested by Kaplan and Nelson \\cite{KN}, and has been discussed in many recent publications \\cite{BLRT,Weise}. Due to the attraction between $K^-$ and nucleons its energy decreases with increasing density, and eventually if it drops below the electron chemical potential in neutron star matter in $\\beta$-equilibrium, a Bose condensate of $K^-$ will appear. It is found that $K^-$'s condense at densities above $\\sim 3-4\\rho_0$, where $\\rho_0=0.16$ fm$^{-3}$ is normal nuclear matter density. This is to be compared to the central density of $\\sim4\\rho_0$ for a neutron star of mass 1.4$M_\\odot$ according to the estimates of Wiringa, Fiks and Fabrocini \\cite{WFF} using realistic models of nuclear forces. The condensate could change the structure and affect maximum masses and cooling rates of neutron stars significantly. Recently, \\cite{PPT,kaon} we have found that the kaon-nucleon and nucleon-nucleon correlations conspire to reduce the $K^-N$ attraction significantly already at rather low densities when the interparticle distance is comparable to the range of the $KN$ interaction. We have calculated the kaon energy as function of density in a simple model where also the low and high density limits and the dependence on the range of the interaction can be extracted. ", + "conclusions": "Kaon-nucleon correlations reduce the $K^-N$ interaction significantly when its range is comparable to or larger than the nucleon-nucleon interparticle spacing. The transition from the Lenz potential at low densities to the Hartree potential at high densities begins to occur already well below nuclear matter densities. For the measured $K^-n$ scattering lengths and reasonable ranges of interactions the attraction is reduced by about a factor of 2-3 in cores of neutron stars. Relativistic effects further reduce the attraction at high densities. Consequently, a kaon condensate is less likely in neutron stars. Coulomb energies have not been included in the above analysis. They may, as discussed in \\cite{Schaffner,rotp}, lead to a mixed phase of nuclear matter with and without a kaon condensate. However, the Coulomb energies are small as compared to kaon masses and therefore the reduction in the kaon energy will also be minor. My collaborators J.Carlson, V.Pandharipande and C.J.Pethick are gratefully acknowledged." + }, + "9802/astro-ph9802254_arXiv.txt": { + "abstract": "We present BeppoSAX observations of the peculiar long period polar system V1309\\,Ori (RXJ0515.6+0105). The source was detected simultaneously at soft and, for the first time, at hard X-rays with the LECS and the MECS detectors. Both, the LECS and the MECS light curves are irregular with a bursting/flaring type behaviour indicating inhomogeneous accretion onto the white dwarf. This peculiar variability, together with an extreme high soft-to-hard X-ray luminosity ratio, indicates that in V1309\\,Ori accretion occurs predominantly in highly compressed chunks or ``blobs'' of matter. From coordinated ESO optical spectroscopy, we find indications that the magnetic field strength of the white dwarf is $<$ 70\\,MG, not expected either from the 8\\,hr orbital period synchronism or from the strong soft-to-hard X-ray ratio suggesting alternative solutions for sustaining synchronism in this system. ", + "introduction": "Polars or AM Her stars are Cataclysmic Variables containing a synchronously rotating magnetic ($\\rm B \\sim$ 10-230 \\,MG) white dwarf accreting from a late-type secondary star. The strong magnetic field of the accreting white dwarf dominates the accretion flow and channels it towards the magnetic polar regions where a strong stand--off shock is produced. The hot post-shock plasma emits hard X-rays, partially absorbed and re-emitted from the surface at soft X-ray, EUV and UV wavelengths, as well as cyclotron emission which is observed at optical and IR wavelengths (Cropper 1990; Beuermann 1997). An independent soft X-ray component can be produced by the infall of dense plasma packets or blobs which penetrate deep into the atmosphere of the white dwarf and heat the photospere from below to a few $10^5$ K (Kuijpers \\& Pringle 1982). Polars have orbital periods in the range from 80\\,min to se\\-ve\\-ral hours, V1309\\,Ori being the longest period system with P$_{\\rm orb}=$ 7.98\\,hr. (Garnavich et al. 1994, Walter et al. 1995, Shafter et al. 1995, Buckley \\& Shafter 1995). V1309\\,Ori is an eclipsing polar showing a deep and variable primary minimum and a shallower secondary one (Garnavich et al. 1994, Shafter et al. 1995, Buckley \\& Shafter 1995). A magnetic field of $\\leq$ 60MG was inferred from cyclotron features in the optical and IR and from optical polarimetry (Garnavich et al. 1994, Shafter et al. 1995, Buckley \\& Shafter 1995, Harrop-Allin et al 1997), although a higher magnetic field was expected for such a long period synchronous system. The soft X-ray emission, as observed by ROSAT, is strongly variable on timescales down to few seconds in a bursting like activity, interpreted by Walter et al. (1995) as evidence for ``blobby'' accretion. The ROSAT data also indicated an X-ray variability at the $\\sim$ 8\\,hr orbital period. Separate ROSAT pointings revealed that V1309\\,Ori shows, as most other polars do, a long-term variability in its X-ray flux due to changes in the mass transfer rate. The soft X-ray emission was found to be consistent with a black-body emission at 50\\,eV. No constraints on the hard component could be established from the ROSAT data. The extraordinary long orbital period of V1309\\,Ori makes it a key object to test theories for synchronization of the white dwarf while its pronounced short-term variability at X-ray wavelengths makes it a test case for the theory of ``blobby'' accretion. In the framework of a program aiming to detect simultaneously the soft and hard X-ray emission in polars with the recent X-ray facility given by the BeppoSAX satellite, we present new X-ray observations of V1309\\,Ori obtained during the BeppoSAX AO1-Core Program together with coordinated optical spectroscopy collected at ESO. ", + "conclusions": "During the BeppoSAX observations V1309\\,Ori was at a flux level comparable to that observed during one of the ROSAT pointings in 1991 when the source was found to be highly variable. A bursting-on/off behaviour, which can be decomposed in quiet intervals interleaved with strong flares, has been observed by BeppoSAX not only in the soft X-rays but also in the hard X-ray emission. Such variability indicates that the activity is due to occasional increases of accretion onto the white dwarf as also suggested by Walter et al. (1995). However the fact that the source sometimes switches off in both hard and soft X-rays is a strong indication of a highly inhomogeneous accretion. Although fla\\-ring acti\\-vi\\-ty has been observed in other polars like BL Hyi (Beuermann \\& Schwope 1989) and QS Tel (Rosen et al. 1996), V1309\\,Ori is the first system to show such a marked variability also in the hard X-rays. Its X-ray spectrum consists of a soft and a hard component which can be described by a 30\\,eV black-body and a 10\\,keV optically thin plasma emission (Raymond-Smith model). This first detection of the hard X-ray component allows us to derive an extremely high soft-to-hard X-ray bolometric flux ratio of $\\sim$ 65-160. Such a large soft X-ray excess indicates that in V1309\\,Ori most of the kinetic energy is emitted from a shock buried deep in the white dwarf atmosphere, at optical depths greater than one even in the hard X-rays. Radiation tran\\-sfer then reprocesses the primary thermal bremsstrahlung into soft X-rays emitted from the surface. We estimate the mean accretion rate assuming that the bulk of the accretion luminosity is irradiated in the soft X-rays: $\\rm L_{bb}=1-2 \\cdot10^{33} erg\\, s^{-1} \\sim G \\dot M\\, M_{wd} R_{wd}^{-1}$. For a standard white dwarf of 0.6$M_{\\odot}$ and a radius of 8$\\cdot10^{8} \\rm cm$, $\\rm \\dot M=1.5-3.2\\cdot 10^{-10} M_{\\odot}\\,yr^{-1}$ which is within the observed range of other polars. A buried shock may form if the white dwarf in V1309\\,Ori possesses a very high magnetic field $\\geq$ 150\\,MG and cyclotron cooling is so efficient that the stand--off shock collapses. Such a high field would be indicated by the 8\\,hr orbital synchronism for a 0.6 M$_{\\odot}$ white dwarf (Patterson 1994), but it is not confirmed by the observations (see below). Another possibility is that the matter is accreted predominantly at high local mass flow rates in form of discrete blobs which penetrate deep in to the atmosphere of the white dwarf (Kuijpers \\& Pringle 1982). Considering that there is no observational evidence for an extremely high magnetic field and that the X-ray light curves display a strong bursting/flaring character, we favour the idea of ``blobby'' accretion to explain the soft X-ray excess in V1309\\,Ori. However, these BeppoSAX observations indicate that the conventional ``blobby '' accretion picture should be modified to account for blobs also to produce hard X-rays. >From optical spectroscopy and polarimetry a magnetic field strength up to 60\\,MG was derived (Garnavich et al. 1994, Buckley $\\&$ Shafter 1995, Shafter et al. 1995, Harrop-Allin et al. 1997). Our optical spectrophotometry acquired during primary minimum and close to the secondary mi\\-ni\\-mum, also suggests a magnetic field of $<$70\\,MG indicating a low field white dwarf, and thus suggesting different solutions to the mechanism of maintaining synchronism. Indeed Frank et al. (1995) proposed that synchronism in V1309\\,Ori, can be sustained if the secondary star possesses a relatively high field ($\\geq$ 1\\,kG). On the other hand a light white dwarf $\\leq 0.47 M_{\\odot}$ is required from the synchronism and for an upper limit of 70\\,MG as also suggested by Shafter et al. (1995). Is then possible that two mechanisms are occurring in V1309\\,Ori: the action of a higher field of the secondary star as proposed by Frank et al. (1995) as well as a lighter accreting white dwarf. \\medskip In summary, while our new X-ray observations indicate that V1309\\,Ori represents an extreme case of ``blobby'' accretion, the optical data support a relatively low field white dwarf thus making this system a test case for further detailed theoretical work." + }, + "9802/astro-ph9802018_arXiv.txt": { + "abstract": "The results of a series of five \\sax observations of the TeV BL Lac object 1ES2344+514 are briefly presented. Large amplitude luminosity variability, associated to impressive spectral changes in the hard X-rays, have been found. The shape of the lightcurve depends on energy, with the flare starting and ending in the hard band, but with maximum intensity possibly reached earlier in the soft X-rays. The luminosity and spectral changes may be due to a shift of the peak of the synchrotron emission from the soft X-rays to the hard X-ray band similar to that detected during \\sax observations of MKN 501. ", + "introduction": "BL Lac objects are a peculiar type of radio loud AGN emitting highly variable non-thermal radiation over an extremely wide energy range from radio waves to TeV energies. Synchrotron emission followed by Inverse Compton scattering is generally thought to be the mechanism responsible for the production of non-thermal radiation over such a wide energy range (e.g. Bregman \\etal 1994). Relativistic beaming is necessary to explain some of the extreme properties of these objects such as rapid variability and superluminal motion (Urry \\& Padovani 1995). 1ES2344+514 is an HBL BL Lac (Padovani \\& Giommi 1995) that is an object where the synchrotron component dominates the spectrum up to very high energies. To date only three BL Lacertae objects (all of them HBLs) have been detected at TeV energies (Catanese et al. 1997): MKN 501, MKN 421 and 1ES2344+514. We briefly present here the main results of a series of broad-band (0.1-200 keV) X-ray observations of 1ES2344+514 carried out with the \\sax satellite (Boella \\etal 1997a). \\begin{figure}[h*] \\epsfig{figure=es2344_pr_fig1.ps, height=7.0cm, width=7.0cm} \\caption{The MECS X-Ray image of the field centered on 1ES2344+514. The bright source at the center is the BL Lac, the faint serendipitous source (1SAXJ2348.7+5128) to the bottom left part of the image coincides with the dwarf nova V630 CAS} \\label{fig1} \\end{figure} ", + "conclusions": "During a series of five \\sax observations 1ES 2344+514 underwent a flare which caused its luminosity to double in the 0.1-10 keV band and to increase of an even larger amount in the PDS band. The X-ray spectral shape of 1ES2344+514 varied with intensity in a way that is typical of HBL BL Lacs, namely the spectrum hardens when the source brightens (e.g. Giommi et al. 1990). The variable X-ray spectral energy distribution of 1ES2344+514 ($\\nu f(\\nu)~vs~\\nu $) is consistent with an interpretation where the flux and the spectral variations are due to the onset of a variable hard component that started dominating the X-ray flux above 5 keV during the second observation and that reached maximum intensity on December 7, when it dominated the entire X-ray band. Figure 4 shows that the peak of the synchrotron power (that is where the spectral slope is 0 in $\\nu f(\\nu)~-~\\nu $ space) during the low state was at frequencies of a few times $10^{17}$ while when the intensity was maximum the peak shifted to $\\approx 3-4\\times 10^{18}$ (20-30 keV). A similar spectral variability, but on an even larger scale, was detected in a recent observation of the X-ray bright and TeV detected HBL BL Lac MKN 501 (Pian et al. 1998). Large shifts of the synchrotron peak energy (implying very large changes in the bolometric luminosity) might therefore be relatively frequent in the hard X-ray band for HBL BL Lacs. Similar behavior should be expected in LBL BL Lacs in the optical/UV band. As for the case of MKN501 (Pian et al. 1998) the results reported here demonstrate that most of the power emitted in BL Lacs may be in the hard X-Rays ($ \\ge 10 $ keV) an energy region poorly explored by previous satellites." + }, + "9802/astro-ph9802197_arXiv.txt": { + "abstract": "Electromagnetic processes associated with a charged particle moving in a strong circular magnetic field are considered in cylindrical coordinates. We investigate the relation between the vacuum curvature emission and Cherenkov emission and argue that, for the superluminal motion of a particle in the inhomogeneous magnetic field in a dielectric, the combined effects of magnetic field inhomogeneity and the presence of a medium give rise to the synergetic Cherenkov-curvature emission process. We find the conditions when the operator relations between electric field and electric displacement in cylindrical coordinates may be approximated by algebraic relations. For nonresonant electromagnetic wave, the interaction with particles streaming along the curved magnetic field may be described in the WKB approximation. For resonant waves, interacting with superluminal particles we use a plane wave approximation to compute the local dielectric tensor of a plasma in a weakly inhomogeneous magnetic field. We find in this approximation the polarization of normal modes in the plasma, Cherenkov-curvature and Cherenkov-drift emissivities and growth rates. ", + "introduction": "Despite much theoretical effort, there is still no widely accepted explanation of how pulsars emit the high brightness radio emission by which they were first discovered over thirty years ago. Difficulties with \"antenna\" mechanisms have led to a resurgence of interest in \"maser\" mechanisms where an inverted population of electrons and/or positrons amplifies an outgoing wave mode (e.g. \\cite{Melrose-DB}). Two such maser process that are particular promising are the Cherenkov-drift mechanism (which is Cherenkov-curvature emission with a drift (\\cite{LyutikovMachabeliBlandford1})) and cyclotron-Cherenkov emission at the anomalous Doppler resonance (\\cite{Kazbegi}). This paper is concerned with elucidating the physics of the Cherenkov-curvature and Cherenkov-drift processes and developing a new mathematical description of it using cylindrical coordinates. Using this approach, its close relationship to the pure Cherenkov emission processes becomes clear. Pulsar radio emission is believed to originate on the open magnetic field lines that trace a path from the neutron star surface to the light cylinder at $r = c/\\Omega$ and beyond. The field geometry is complex close to the star at $ r \\approx R_{\\ast} \\approx 10 $ km and near the light cylinder. However, it should be primarily dipolar for $R_{\\ast} \\ll r \\ll c/\\Omega$. The characteristic length scale, essentially the radius of curvature of the field lines, $ R_B \\approx (c r /\\Omega)^{1/2}$, is very large compared with the wavelength of interest, $\\lambda$, and so we should be able to adopt a WKB approach following following individual wave packets as they propagate out through the magnetosphere. However, in order to compute the local emission and absorption it is necessary that the magnetic field remains curved. A simple model problem, that retains the essential ingredients, comprises a set of circular concentric cylindrical magnetic surfaces. (Fig. \\ref{max1}). We show, that it is possible to ignore the radial variations in the strength of ${\\bf B}$ in computing the local interaction, provided that this variation is not very rapid. In this model problem the plasma circulates continually along the circular trajectories. It is generally presumed, that electrons and positrons are in their ground gyrational state and follow the curved field lines. This is a good approximation for the bulk of the plasma, which we will suppose travels with a Lorentz factor $\\gamma_p \\approx 100$. However, we believe that there is also a population of ultra energetic particles with $\\gamma =10^{5-7}$ and, in the outer magnetosphere, these will experience a curvature drift relative to the bulk plasma. This has the interesting consequence that their main electrodynamic interaction is with the waves propagating in the bulk plasma at a finite angle to the magnetic field ${\\bf B}$. It is these waves, that, we assert, become the high brightness electromagnetic waves that escape from the magnetosphere. When considering the wave amplification in this problem, we assume that fluctuating currents, present on the field lines at radii smaller than the radius of the considered region, produce electromagnetic fluctuations that are subsequently amplified due to the interaction with resonant particles (Fig. \\ref{max1}). This is a convective type instability, when a wave is amplified as it propagates through an active medium. A first attack on this problem was made by \\cite{Blandford1975} In this paper the energy transfer from a plane infinite electromagnetic wave to a single electron moving with ultrarelativistic speed along a curved trajectory was found to be always positive, independent of the distribution function and so there is no possibility of wave growth. The subsequent works (\\cite{ZheleznyakovShaposhnikov}, \\cite{MelroseLou}, \\cite{Melrosebook1}) basically followed this approach, which emphasizes the analogy between curvature emission and conventional cyclotron emission. This approach, though formally correct, has limited applicability and ignores two important features of the emission mechanism. The first is that, in adopting a plane wave formalism, the interaction length for an individual electron, $\\approx R_B /\\gamma_b$, was essentially coextensive with the region over which the waves could interact with any electron. This approach precludes a strong amplification under all circumstances because the wave would have to grow substantially during a single interaction. The second problem was that a dispersion of the waves was neglected. We address the first shortcoming by expanding the electromagnetic field in cylindrical waves centered on $r=0$, and the second explicitly by considering general plasma modes. In a separate approach developed by \\cite{beskin1} an attempt was made to incorporate the collective effects of a plasma in an inhomogeneous magnetic field. The fundamentals of that approach have been seriously criticized (\\cite{Nambu}, \\cite{machabeli}). One of the key new elements in our approach is that we expand the electromagnetic fields in cylindrical coordinates and consider {\\it resonant } interaction between a particle and these modes. A common procedure in calculating the energy emitted by a given current distribution is to find the power emitted into a given normal mode of the medium and sum over all the modes. A power emitted into a normal mode by a given charge distribution is proportional to the square of the expansion of this current in normal modes. In a homogeneous or weakly inhomogeneous stratified medium considered in Cartesian coordinates the normal modes are plane waves (with slowly changing parameters), so that the power emitted into a given normal mode turns out to be proportional to the Fourier transform of the current. To find a similar expression for a single particle emissivity into a cylindrical mode in an inhomogeneous medium is difficult problem for two reasons: there is a complicated, radius dependent, form of the vector cylindrical waves and the dielectric response in an inhomogeneous medium is nonlocal. We address these complications by finding the conditions, when the interaction of a particle with cylindrical waves can be approximated by an interaction of a particle with plane waves. We find two cases when this can be done. First, this may be a true approximation for the{\\it nonresonant} modes. This is equivalent to the WKB approximation to the radial dependence of normal modes (mathematically, this corresponds to the tangent expansion of Bessel functions when argument is significantly larger that he order). In this case the response of a medium becomes local. Second, we find a particular case, when the {\\it resonant} modes can be approximated as local plane waves. A resonant interaction of a relativistic particle with a cylindrical mode occurs near the point when the argument of Bessel functions is close to the order. The WKB approximation, or expansion in tangents, is not applicable in this case and we have to use the Airy function approximation to Bessel function, which has a plane wave approximation for the interaction of {\\it subluminous} waves with the particles moving with speed larger than the speed of light in a medium This corresponds to the Airy function expansion argument being larger than the order. We should note here, that the electrodynamics of the interaction of a particle moving along the curved magnetic field with speed larger than the speed of light in the medium is quite unusual and can be considered as a new, Cherenkov-Curvature, emission mechanism (\\cite{LyutikovMachabeliBlandford1}), which differs from conventional Cherenkov, cyclotron or curvature emission and includes, to some extent, the features of each of these mechanisms. In an extension of ideas of \\cite{Schwinger}, conventional synchrotron emission and Cherenkov radiation may be regarded as respective limiting cases of $|n-1|\\, \\rightarrow \\,1 $ and $B \\rightarrow \\,0$ of a \"synergetic\" cyclotron-Cherenkov radiation, the Cherenkov-curvature radiation necessary includes effects of the magnetic field gradients. When the dielectric response of a medium to cylindrical wave is {\\it local} it is possible to calculate a simplified dielectric tensor. Using this dielectric tensor we find the normal modes of strongly magnetized electron-positron plasma and show that a beam of particles propagating along the curved magnetic field can amplify the electromagnetic waves. The amplification occurs at the Cherenkov-drift resonance $\\omega-k_{\\phi} v_{\\phi} - k_x u_d=0$ ($\\omega$ is the frequency of the wave, $k_{\\phi}$, $k_x$ and $v_{\\phi}$ are the corresponding projection of the wave vector and velocity and $u_d$ is a curvature drift velocity). The fact that this is a Cherenkov-type resonance immediately implies that the presence of a subluminous waves with the phase velocity smaller that the speed of light is essential. The presence of a drift provides a coupling between the the electric field of the electromagnetic waves particle's motion. Our estimates show, that this instability can grow fast enough to account for the observed pulsar radio emission. This instability may be regarded as new type of a curvature maser. An interesting and peculiar feature of this mechanism, is that, unlike with conventional curvature emission, the emitted waves have a polarization almost perpendicular to the osculating plane of the magnetic field. The overview of this work is the following. In Section \\ref{vacum} we analyze the properties of electromagnetic waves in cylindrical coordinates in vacuum and the interaction of the vacuum waves with a charged particle. In Section \\ref{SingleParticle} we analyze in cylindrical coordinates the electromagnetic fields of a particle moving along a spiral trajectory using the dyadic Green's function for the vector wave equation, derive the emissivity of a particle in ground gyrational level into a cylindrical mode and rederive the curvature emissivity in cylindrical coordinates. This is followed in Section \\ref{Wavesisotroipicdielectric} by the generalization to a dispersive isotropic medium and the importance of waves with the phase speed less than the speed of light in vacuum is brought out. In Section \\ref{DispersionInf} we consider {\\it anisotropic} plasma in infinitely strong magnetic field and derive the curvature emissivity using the Vlasov approach. In Section \\ref{Airyfunction} we discuss the various regimes of the Airy function approximation to the Cherenkov-curvature emission and find the conditions for the plane-wave approximation. In Section \\ref{WavesAsymptotic} we investigate the features of the electromagnetic waves in the asymptotic regime $z \\gg 1$ and in Section \\ref{ResponceTensor} we calculate the response tensor for a one dimensional plasma in a strong curved magnetic field taking into account the drift velocity. Finally, in Section \\ref{pl} we analyze the polarization properties of electromagnetic waves in the plane-wave approximation and calculate the increment of the Cherenkov-drift instability. ", + "conclusions": "\\label{concluss} { In this work we investigated electromagnetic processes associated with a charged particle moving in a dielectric in a strong circular magnetic field. We derived a simple expression for the growth rate of the Cherenkov-drift instability that, we believe, may be responsible for the generation of the pulsar radio emission. We leave the discussion of the astrophysical applications of our results for the subsequent paper (\\cite{LyutikovBlandfordMachabeli}). Here we just note, that the Cherenkov-drift developing on the open field lines in the outer parts of pulsar magnetosphere can explain various features of the cone-type (\\cite{rankin1}) emission patterns observed in pulsars. }" + }, + "9802/astro-ph9802183_arXiv.txt": { + "abstract": "We present a combined X-ray and optical analysis of the ABCG 85/87/89 complex of clusters of galaxies, based on the ROSAT PSPC image, optical photometric catalogues (Slezak et al. 1998), and an optical redshift catalogue (Durret et al. 1998). From this combined data set, we find striking alignments at all scales at PA$\\simeq$160\\deg. At small scales, the cD galaxy in ABCG 85 and the brightest galaxies in the cluster are aligned along this PA. At a larger scale, X-ray emission defines a comparable PA south-southeast of ABCG 85 towards ABCG 87, with a patchy X-ray structure very different from the regular shape of the optical galaxy distribution in ABCG 87. The galaxy velocities in the ABCG 87 region show the existence of subgroups, which all have an X-ray counterpart, and seem to be falling onto ABCG 85 along a filament almost perpendicular to the plane of the sky. To the west of ABCG 85, ABCG 89 appears as a significant galaxy density enhancement, but is barely detected at X-ray wavelengths. The galaxy velocities reveal that in fact this is not a cluster but two groups with very different velocities superimposed along the line of sight. These two groups appear to be located in intersecting sheets on opposite sides of a large bubble. These data and their interpretation reinforce the cosmological scenario in which matter, including galaxies, groups and gas, falls onto the cluster along a filament. ", + "introduction": "One of the major changes in our understanding of our universe has been the realization of the rich and complex structures which are apparent in the large scale distribution of galaxies. From both redshift surveys and projected galaxy distributions, the appearance of galaxy voids and supercluster filaments has become clear. The large scale structure also is an important constraint for different cosmological scenarios. On smaller scales, investigations of the relationships between nearby galaxy clusters suggest that clusters retain information about the large scale structures from which they form (e.g., van Haarlem \\& van de Weygaert 1993; West et al. 1995, West 1997; Colberg et al. 1997). The frequency of substructure may also provide constraints on $\\Omega$ (Richstone et al. 1992, Mohr et al. 1995, Buote \\& Xu 1997; see also Kauffmann \\& White 1993). As part of a survey for substructure, we have studied the large scale mass distribution around the richness class~1 cluster ABCG 85 (Abell et al. 1989), using X-ray and optical observations. As a bright, luminous, relatively nearby cluster, ABCG 85 has been studied extensively. At a redshift of $z=0.0555$, the angular scale for ABCG 85 corresponds to $72~h_{50}^{-1}$~kpc arcmin$^{-1}$ (we assume H$_0$= 50 km sec$^{-1}$ Mpc$^{-1}$ throughout). ABCG 85 contains a cD galaxy close to its center. The cluster X-ray characteristics include a peaked emission profile, harboring a cooling flow, emission from individual cluster member galaxies, an X-ray emitting subcluster south of the cluster center (hereafter the south blob), a superposed foreground group of galaxies (to the west-northwest), as well as additional foreground and background structures detected from the optical spectroscopic observations (see Jones et al. 1979, Pislar et al. 1997, Lima-Neto et al. 1997, Slezak et al. 1998, Durret et al. 1998, and in preparation for detailed discussions of the X-ray and optical observations). In this contribution, we report on the X-ray and optical properties of a larger region around ABCG 85 than has been considered in previous studies, and in particular include the nearby clusters ABCG 87 and ABCG 89. We find that ABCG 85 itself exhibits preferential alignments on scales from 100 kpc to $\\sim 4$Mpc (in projection on the sky). In particular, the cluster cD galaxy is elongated along the same position angle as a large filamentary structure in X-rays roughly coinciding with the optical position of ABCG 87. Located to the east of ABCG 85 is ABCG 89, not detected as an extended X-ray source. We also address why ABCG 87 appears different in the optical and X-ray and why ABCG 89 appears as a cluster in the optical observations and not in X-rays. Based on these results, we discuss the ABCG 85/87/89 complex and describe the alignments and/or structures in the context of a large scale structure formation scenario such as that proposed by West et al. (1995; see also van Haarlem \\& van de Weygaert 1993 and Colberg et al. 1997). In Sect. 2 we present the X-ray imaging analysis and galaxy distributions. Sect. 3 discusses in detail the galaxy velocity distribution. A model is proposed in Sect. 4 and discussed in Sect. 5. \\begin{figure}[tbp] \\centerline{ \\psfig{figure=h0823.f1,height=8cm}} \\caption{Optical digitized sky survey of the ABCG 85 region, with the isophotes of the ROSAT PSPC image of a flat-fielded $55'$ radius field superimposed. The energy band used is 0.4-2.0 keV and the image has been smoothed with a $60''$ Gaussian after flat-fielding. The cD galaxy at the X-ray cluster center (peak of the X-ray emission) is clearly visible on the optical map; a bright star is located in the region of extended X-ray emission, but it does not significantly contribute to the observed X-ray emission.} \\protect\\label{xcont} \\end{figure} ", + "conclusions": "We presented a three dimensional model of the complex ABCG 85/87/89 region. We find a general extension observed toward the south-east along PA$\\simeq 160$\\deg\\ both in X-rays and in optical photometric observations: - In X-rays, the elongated structure extends to the edge of the ROSAT PSPC field of view; the X-ray appearance includes several individual group-like systems lying up to about 4~Mpc (in projection) from the center of the cluster (see Fig.~\\ref{xcont}) as well as very elongated diffuse emission; - from the optical photometric catalogue, ABCG 85 is elongated towards the ABCG 87 cluster (see Fig.~\\ref{kernel}). The PA$\\simeq 160$\\deg\\ is characteristic for ABCG 85: the direction of the major axis, the brightest galaxies in ABCG 85, and even the major axis of the central cD galaxy itself are elongated along the same direction. East of ABCG 85 lies ABCG 89 which is not bright in X-rays. Our velocity data reveal that ABCG 89 is not a cluster, but the superposition on the sky of two groups well separated in velocity space. These two groups (to which we refer as A89b \\& A89c in Fig.~\\ref{artist}) are located in intersecting sheets on opposite sides of a large bubble. We have shown that ABCG 87 is not a rich cluster, but is resolved into individual groups possibly falling onto ABCG 85. These groups are organized as a filament almost perpendicular to the plane of the sky. The superposition of these groups gives the appearance of a single optical cluster, while in X-rays, several groups are still visible: This is probably due to the fact that the emissivity in X-rays is proportional to the density squared, therefore enhancing the contrast, compared to optical imaging. ABCG 85 itself is probably not fully relaxed, even if it appears smooth and symmetric; the distribution of velocities is obviously not gaussian and probably multi-modal. In his scenario for the Coma cluster, West (1997) links the orientations of Coma to the filament in which Coma is embedded. He suggests that matter, including galaxies, groups and gas, falls onto the cluster along this filament. The case of ABCG 85 reinforces this cosmological scenario because \\textit{we actually observe} the infall of material (a filament of groups of galaxies and gas) onto ABCG 85. This result is consistent with the X-ray temperature map derived from ASCA data by Markevitch et al. (1998), which shows a temperature enhancement in a region south of ABCG 85 and roughly perpendicular to the general direction of 160\\deg\\ along which the various structures are aligned. Such a temperature increase could be interpreted as shock heating due to the compression of the X-ray gas by infalling matter. The fact that there also appears to be a radio relic, also roughly perpendicular to the 160\\deg\\ direction, in this zone (Bagchi et al. 1998) also supports recent merger activity, since relativisitic electrons can be produced during a merger. Remarkably, the ABCG 85/87 filament is coaligned with a much larger structure including from northwest to southeast: ABCG 70, ABCG 85 and 89, ABCG 87, ABCG 91, the NGC 255 group and ABCG 106 (Fig.~1 in Slezak et al. 1998). Such a structure extends over more than 5\\deg\\ on the sky, corresponding to a linear distance of 28 Mpc at the redshift of ABCG 85 (z=0.0555). This is obviously a lower limit, since the filament may be inclined to the line of sight. Although this projected value is smaller than that found for example for the Perseus Pisces structure (50$h^{-1}$ Mpc, Haynes \\& Giovanelli 1986), it is nevertheless much larger than typical cluster sizes. Unfortunately, there are only very few redshifts available for this large structure outside the ABCG 85/87/89 complex. Further optical observations are required to confirm the cohesiveness of this remarkable filamentary structure." + }, + "9802/astro-ph9802240_arXiv.txt": { + "abstract": "This is the third in a series of papers in which spiral galaxy rotation curves are considered in the context of Milgrom's modified dynamics (MOND). The present sample of 30 objects is drawn from a complete sample of galaxies in the Ursa Major cluster with photometric data by Tully et al. (1996) and 21 cm line data by Verheijen (1997). The galaxies are roughly all at the same distance (15 to 16 Mpc). The radio observations are made with the Westerbork Synthesis Array which means that the linear resolution of all rotation curves is comparable. The greatest advantage of this sample is the existance of K$^{'}$-band surface photometry for all galaxies; the near-infrared emission, being relatively free of the effects of dust absorption and less sensitive to recent star formation, is a more precise tracer of the mean radial distribution of the dominant stellar population. The predicted rotation curves are calculated from the K$^{'}$-band surface photometry and the observed distribution of neutral hydrogen using the simple MOND prescription where the one adjustable parameter is the mass of the stellar disk or the implied mass-to-light ratio. The predicted rotation curves generally agree with the observed curves and the mean M/L in the near-infrared is about 0.9 with a small dispersion. The fitted M/L in the B-band is correlated with B-V color in the sense predicted by population synthesis models. Including earlier work, about 80 galaxy rotation curves are now well-reproduced from the observed distribution of detectable matter using the MOND formula to calculate the gravitational acceleration; this lends considerable observational support to Milgrom's unconventional hypothesis. ", + "introduction": "The rotation curves of spiral galaxies, as observed in the 21 cm line of neutral hydrogen, provide, in many cases, an accurate determination of the radial force distribution in regions of very low gravitational acceleration ($<10^{-8}$ cm/s$^2$). Therefore, these data are ideal for testing alternatives to dark matter such as Milgrom's proposed modified Newtonian dynamics (MOND) which posits that the discrepancy between the true gravitational force and the Newtonian force appears at low accelerations. As a theory of gravity or inertia, MOND predicts the precise form and amplitude of a rotation curve from the observed radial distribution of detectable matter (stars and gas) in a spiral galaxy, often with only one adjustable parameter which is the mass-to-light ratio of the stellar disk (Milgrom 1983a, b). In two previous papers (Begeman et al. 1990 Paper 1, Sanders 1996 Paper 2), a sample of 33 spiral galaxies with published rotation curves has been considered in the context of MOND. The 11 galaxies considered in Paper 1 were highly selected to meet certain strict criteria: the galaxies were rich in neutral hydrogen; the distribution and velocity field of the neutral gas was smooth and symmetric; the observed 21 cm line rotation curves extended far beyond the visible disk; the galaxies were closer than 1000 km/s to achieve high linear resolution and a large number of independent points along the rotation curve. In Paper 2 these criteria were relaxed to include an additional 22 galaxies with published 21 cm line rotation curves (as of 1996), but selection was still essentially based upon HI richness and extent. For the combined samples of Papers 1 and 2 it was demonstrated that, in most cases, the observed rotation curve was predicted in detail from the observed light and gas distributions using the simple MOND prescription. Moreover, the range of fitted values of M/L was astrophysically plausible and consistent with population synthesis models. In the present paper, this work is extended to a new complete and homogeneous sample of spiral galaxies in the Ursa Major cluster. The sample has been previously considered in terms of the distribution by surface brightness (Tully \\& Verheijen 1997), but until now, the dynamics of these galaxies has not been discussed in the published literature. The important aspect of this sample is that it is optically selected (Tully et al. 1996). All galaxies in the UMa cluster brighter than a specified limiting magnitude are considered; although, the final sample is weighted against gas-poor systems because the rotation curves are determined from 21 cm line observations made at the Westerbork Radio Synthesis Telescope (Verheijen 1997). All sample galaxies have been imaged in the K$^{'}$-band (2.2 $\\mu$m) by Tully et al. (1996). Therefore, this sample is distinguished from the previous samples in its primary selection criterion and by its homogeneity: rotation curves of consistent and sufficient linear resolution are combined with near-infrared surface photometry in a data set which is well-suited to the purpose of this work. The previous combined sample of 33 field galaxies (Paper 2) was very inhomogeneous with the 21 cm line observations being carried out at either Westerbork or the VLA with differing angular resolution and sensitivity. Considering the range of distance covered by the sample (from 0.8 to 80 Mpc) the actual linear resolution of the neutral hydrogen observations varied between about 50 pc and 5 kpc. The 21 cm line data were taken by different observers using different methods of analysis. The photometry was quite inhomogeneous; in most cases this was CCD photometry but for several objects only older photographic photometry was available. The photometry was typically, but not always, in the B-band which is less reliable than redder bands as a tracer of the true radial distribution of the dominant stellar population, subject as it is to recent star-formation and the effects of differential dust obscuration. Moreover, nine of these galaxies contained bulges or central light concentrations which, at least in the B-band, yielded yet another fitting parameter-- the M/L of the bulge component. These problems are minimized in the UMa sample. All galaxies are roughly at the same distance (taken to be 15.5 Mpc) which means that the linear resolution of the 21 cm line observations is similar in all cases (typically 0.75 kpc). The relatively low scatter in the observed Tully-Fisher relation implies a rather small dispersion in distance. The observations are reduced in the same way and the same method is applied in deriving the rotation curves from the line data (Verheijen 1997). One of the greatest advantages of this sample is that the radial distribution of the old stellar population is determined from near-infrared (K$^{'}$-band) CCD photometry which is free from the above-mentioned problems associated with the B-band. The galaxies are generally of quite late morphological type and the effects of a separate bulge component with a spheroidal shape and separate M/L are minimal. Indeed, the M/L in the near-infrared is probably similar for the disk and the presumably older bulge component. It is important to bear in mind that this is an {\\it optically} selected sample; these galaxies are not selected for HI richness or the large extent of the neutral hydrogen disks. In many of these objects the neutral hydrogen does not extend far beyond the optical disk into the region of low accelerations where the discrepancy between the visible and classical dynamical mass is large. Both high and low surface-brightness galaxies are present, as well as a range of morphological types. Overall, this must be considered an advantage; it cannot be argued that these galaxies were picked as cases particularly favorable to analysis in terms of MOND. Moreover, the range of objects and accelerations probed illustrate some significant general characteristics of the mass discrepancy in spiral galaxies. However, this complete and homogeneous sample is not completely free from problems. Because it is an optically selected sample, in a number of these objects the HI mass is low and its distribution is patchy-- not ideal for derivation of a high-quality rotation curve. The distance to these galaxies is rather large compared to that of the galaxies in the highly selected sub-sample of Paper 1. Combined with the fact that a number of the galaxies have a small angular size, the number of independent points on the measured rotation curve is, in several cases, quite low (less than five). The relatively poor linear resolution means that real structure in the rotation curve can be artificially smoothed (beam smearing). These galaxies are members of a cluster, albeit a loose cluster, and therefore quite a number of objects are interacting with near neighbors. This may complicate the interpretation of the neutral hydrogen velocity field as pure circular motion about the center of the galaxy. Taking these factors into account, the measured rotation curves, as tracers of the radial force law, are generally inferior to those of Paper 1, although the photometry is superior and the relative error in distance is smaller. With these caveats in mind, we have carried out the MOND analysis on the galaxies in the UMa sample. This adds between 20 and 30 rotation curves (depending upon internal selection criteria) to the total number considered in the context of this theory. The results are generally positive for modified dynamics; the predicted rotation curves agree with the observed curves and the range of implied M/L values is reasonable. In the K$^{'}$ band, the scatter in M/L is strikingly small: the mean M/L is about one in solar units with a dispersion of 30\\%. This is entirely consistent with the scatter in the observed K$^{'}$-band TF relation (Verheijen 1997); in the context of MOND the only source of intrinsic scatter in the TF relation is that in M/L, assuming planar gas disks in circular rotation. Below we describe the sample, show the predicted and observed rotation curves and discuss the implications for the hypotheses of modified dynamics and dark matter. ", + "conclusions": "This sample of galaxies in the Ursa Major cluster combined with those of Paper 2 and the low surface-brightness galaxies analyzed by McGaugh and de Blok (1998b), provide a total sample of about 80 spiral galaxies with rotation curves measured in neutral hydrogen for which the predicted MOND rotation curves have been calculated. Because of the systematic effects the usual statistical tests of goodness-of-fit are not entirely appropriate in accessing these rotation curve fits, but qualitatively one would say that in only five or six cases out of this 80 is the rotation curve predicted by MOND from the observed distribution of light and gas noticeably different from the observed rotation curve. In the present sample, NGC 4389 and UGC 6973 should be included in this category. For these objects there is usually an obvious problem with the observed rotation curve or its use as a tracer of the radial force distribution. In others one may question the measured surface brightness distribution as a precise tracer of the stellar mass distribution (UGC 6973 for example). For about 30 galaxies the MOND rotation curve reproduces the general shape of the observed curve and accounts for the magnitude of the conventional mass discrepancy with a reasonable mass-to-light ratio (NGC 3877, 4013, 4138 would fall in this category). In fully half of the objects considered, the MOND rotation curve reproduces the observed curve with considerable precision (as for NGC 4157, UGC 6399, 6667, 6917 in the present sample). Unlike the previous samples, the galaxies considered here are optically selected and form a complete and homogeneous sample. Because selection does not depend primarily upon neutral hydrogen richness and extent, these galaxies have a wider range of properties than those in Papers 1 and 2. For example, in a number of these objects the observed rotation curve does not extend far beyond the optical disk into the regime where the discrepancy is large. This is a positive aspect of this sample because certain general predictions of MOND are seen to be verified: the discrepancy between the Newtonian and observed rotation curve is seen to be small in those objects with high centripetal accelerations at the last measured point of the rotation curve-- as MOND predicts. In regions where the internal acceleration is low-- in the extended rotation curves of high surface-brightness galaxies as in the previous samples-- MOND predicts the observed large discrepancy. It is noteworthy that some of the best fits are achieved for the low surface-brightness galaxies where the discrepancy is largest. This is precisely what one would expect from MOND since these objects lie entirely in the low acceleration limit of the theory-- in the deep MOND limit-- where there is no uncertainty arising from unknown function $\\mu(x)$ in eq.\\ 2 (see also McGaugh \\& de Blok 1998b). The range of M/L is generally reasonable and consistent, in the blue, with population synthesis models. When near-infrared photometry is available, as for this sample, there is remarkably little variation in the fitted M/L of the stellar component-- on the order of 30\\%. When one considers that this is the only free parameter in the fits and will reflect all uncertainties (such as the real dispersion in distance), the results here are consistent with a effectively constant value of M/L in the near-infrared for spiral galaxies. Assigning the same near-infrared M/L to all galaxies then makes the MOND rotation curves true predictions, without free parameters. The predictive power of MOND, at least with respect to galaxy rotation curves and the TF relation, is well-established. It has been difficult to compare this, in a fair way, with dark matter halos because, until recently, the dark halo hypotheses have had essentially no predictive power. The exercise has been one of fitting dark matter halos with an assumed density distribution to the observed rotation curve in order to estimate halo and disk parameters. With at least three free parameters available (disk M/L, halo core radius, and density normalization), essentially any observed rotation curve can be reproduced, as in the analysis of UMa galaxies by Verheijen (1997). In the last several years, this has changed due to the further development of cosmological N-body codes with large numbers of particles and high spatial resolution. It has become evident that in simulations in which the initial spectrum of density fluctuations is that of Cold Dark Matter (CDM), a characteristic form for the radial mass distribution of virialized halos arises (Dubinsky \\& Carberg 1991, Navarro et al. 1996, Cole \\& Lacey 1996). This characteristic form is distinguished by a singular density distribution-- a density which continues to increase as 1/r into the origin. This density law may be parameterized by the Hernquist model (Hernquist 1990, Dubinsky \\& Carlberg 1991) or by an alternative model suggested by Navarro et al. (1996) which has essentially the same form in the regime where galaxy rotation curves actually measured. Although these singular halos can produce acceptable fits to the rotation curves of high luminosity, high surface-brightness galaxies (Sanders and Begeman 1994) they generally fail in low surface-brightness dwarfs (Flores \\& Primack 1994, Moore 1994). Significantly, Navarro et al. (1996) have demonstrated that, given the mass or velocity scaling of a halo, the degree of central concentration (or halo length scale) is determined by the cosmology. McGaugh \\& de Blok (1998b) show that such constrained halos dramatically fail to reproduce the rotation curves of the low surface-brightness galaxies for plausible cosmologies; it is impossible to match the gradual rise in the observed rotation curve in the inner region with the asymptotic velocity at larger radii; the singular density distribution produces a rotation curve which rises far too steeply. The predicted mass-velocity relation for the CDM halos (not considering that the observed correlation is with the luminosity of the baryonic component) is closer to $M\\propto V^3$ (White 1997) while $M\\propto V^4$ for MOND-- again a clearly distinct prediction. In this sense also, MOND is successful when one considers the correlation between the asymptotic flat rotation velocity and the luminosity (Verheijen 1997). Therefore, on this very basic phenomenological level, one can only conclude that MOND works where dark matter, at least the presently favored form of dark matter, does not. It is sometimes argued that MOND is ``designed'' to fit rotation curves, so that it is no surprise that it works so well on this scale. It is true that MOND, in some sense, is designed to reproduce asymptotically flat rotation curves and a TF relation of the form $L\\propto V^4$. But it by no means evident that the variety of detailed shapes of rotation curves exhibited by the total sample of 80 galaxies so far considered could be so precisely reproduced by using MOND to calculate the radial force from the observed distribution of detectable matter. MOND works well throughout the entire galaxy, not just where the rotation attains its asymptotic constant value, and it was in no sense designed to do this. Using MOND Milgrom {\\it predicted} that the discrepancy between the observed curve and the Newtonian rotation curve should be small in regions of high surface-brightness, and that the discrepancy should be large in galaxies of low surface-brightness-- even before such galaxies were discovered. These predictions are clearly verified in the sample of galaxies considered here, which include both types of systems. It is this inherent simplicity and predictive power that gives the idea its continued impact in spite of the absence of a solid theoretical basis for MOND. The detailed shape and amplitude of a rotation curve can be calculated from the observed distribution of detectable matter, with no free parameters by assuming a fixed value of M/L of about one in the near-infrared. What more could one expect from a theory of gravity? Because of this it is justified to claim that MOND, at the present time, is epistemologically superior to the dark matter hypothesis." + }, + "9802/astro-ph9802289_arXiv.txt": { + "abstract": "We present ISO-SWS spectroscopy of the cool dusty envelopes surrounding two Planetary Nebulae with [WC] central stars, BD+30~3639 and He~2-113. The $\\lambda$~$<$~15~$\\mu$m region is dominated by a rising continuum with prominent emission from C-rich dust (PAHs), while the long wavelength part shows narrow solid state features from crystalline silicates. This demonstrates that the chemical composition of both stars changed very recently (less than 1000 years ago). The most likely explanation is a thermal pulse at the very end of the AGB or shortly after the AGB. The H-rich nature of the C-rich dust suggests that the change to C-rich chemistry did not remove all H. The present-day H-poor [WC] nature of the central star may be due to extensive mass loss and mixing following the late thermal pulse. ", + "introduction": "Planetary Nebulae (PNe) are the ionized and photodissociated remnants of extensive mass loss which the central star experienced when it was a cool Asymptotic Giant Branch (AGB) star. PNe therefore provide insight in the mass loss on the AGB and beyond. An interesting class of PNe are those with a Wolf-Rayet ([WC]) central star. These central stars have strongly enhanced C and He, but little or no H in their atmosphere, and are the low mass counterparts to the population~I Wolf-Rayet stars. About 50 [WC] stars are known (Gorny \\& Stasinska 1995). The properties of these nebulae do not differ significantly from nebulae with 'normal' central stars (Pottasch 1996). The formation of H-poor central stars is somewhat of a puzzle, but is probably related to a thermal pulse either at the very end of the AGB or young PN phase (in this paper referred to as PAGB pulse; Zijlstra et al. 1991), or when the star is already on the cooling track (very late thermal pulse; Iben 1984). The H-rich layers may be removed due to efficient mixing and subsequent nuclear burning, or by extensive mass loss, exposing processed layers to the surface. Examples of objects that may have recently experienced a late thermal pulse are FG~SGe and Sakurai's object (e.g. Paczy\\'nski 1970; Nakano et al 1996; Duerbeck et al. 1996; Bl\\\"ocker \\& Sch\\\"onberner 1997). It is not clear that these stars will develop [WC] type spectra. IR spectroscopy of the dust ejected by [WC] central stars shows the well-known family of IR emission bands, usually identified as due to Polycyclic Aromatic Hydrocarbons (PAHs) (e.g. Cohen et al. 1986). This confirms the C-rich nature of the most recent mass loss episode. In this {\\em Letter}, we present new infrared spectra taken with the Short Wavelength Spectrometer (SWS) on board of the Infrared Space Observatory (ISO) of two well-studied PNe with [WC] central stars, BD+30~3639 and He~2-113. Both nebulae have associated molecular gas (Taylor et al. 1990; Bachiller et al. 1991; Gussie \\& Taylor 1995). In Sect.~2, we discuss the observations and data reduction, Sect.~3 gives an inventory of the solid state features (both C-rich and O-rich). In Sect.~4 we present a first attempt to model the dust spectrum of BD+30~3639, and we discuss these results in the context of post-AGB evolution and formation scenarios for [WC] central stars of PNe. \\begin{figure} \\psfig{figure=fig1.ps,width=10cm} \\caption[spectra]{2.4-45 $\\mu$m ISO-SWS spectra of He~2-113 and BD+30~3639, showing the PAH emission features at $\\lambda$ $<$ 13 $\\mu$m and the crystalline silicates at $\\lambda$ $>$ 20 $\\mu$m. Tickmarks indicate olivine bands (short thin lines; Dorschner, private communication) and pyroxene (long thick lines; Koike et al. 1993). Emission from [Ne~II] at 12.8 $\\mu$m is seen in both nebulae, and from [Ar~II] at 6.99 $\\mu$m, [S~III] at 18.71 $\\mu$m, and [Si~II] at 34.8 $\\mu$m in BD+30~3639.} \\end{figure} ", + "conclusions": "\\subsection{A dust model} The presented model is calculated using the dust radiation transfer code {\\sc modust} (de Koter et al. in preparation). We have modeled the SWS spectrum of BD+30~3639 using a two component dust model, of which the inner and outer shell connect. The inner part contains amorphous carbon grains and we have constrained the radial dimensions of this shell by requiring that it coincides with the region from which PAH emission is observed (Bernard et~al. 1994), i.e. from about 1.9'' to 2.7''. Adopting a distance of 2.6 kpc (Hajian \\& Terzian 1994), this yields an inner and outer radius of the C-rich shell of 7.3~10$^{16}$ and 10.6~10$^{16}$ cm respectively. The luminosity of BD+30~3639 was set to 22,000 L$_{\\odot}$ and its effective temperature to 34.2 kK. Both these values are derived from the analysis of Siebenmorgen et~al (1994) after scaling to our (larger) adopted distance. Note that T$_{\\rm eff}$ is rather uncertain as estimates of this value range from 30 kK (Siebenmorgen et~al. 1994) to 47 kK (Leuenhagen et~al. 1996). The stellar luminosity yields a current mass $M$ = 0.89 $M_{\\odot}$ (Paczy\\'nski 1970), neglecting the possible effects of hot-bottom burning. Radiative transfer in the inner dust region is properly taken into account and the emerging spectrum is used to irradiate the outer shell. The outer shell is assumed to be optically thin. Feed-back from the outer to the inner shell has not been taken into account, which is a valid assumption as the O-rich outer shell is optically thin at all wavelengths. The outer shell contains a mixture of amorphous and crystalline silicates. We use olivine (Dorschner, private communication), ortho-pyroxene and clino-pyroxene (Koike, private communication), and amorphous silicates (Draine \\& Lee 1984) with 12, 1, 2 and 85 per cent abundance respectively (by mass), and a grain size radius of 0.2 $\\mu$m (0.25 $\\mu$m for the amorphous silicates). The outer radius of the O-rich shell is 4.4~10$^{17}$ cm, but is poorly constrained by the model. The resulting model fit is shown in Figure~3. Note that our solid state model does not yet include all structures seen in the spectrum; this will improve with better laboratory data. We find mass-loss rates for the C-rich and O-rich shells of 2~10$^{-5}$ and 8~10$^{-5}$ M$_{\\odot}$/yr respectively, assuming a gas over dust ratio of 100 in both shells. The errors in the mass loss are at least a factor of two. Reasons for this large uncertainty are, among others: {\\em (i)} the poorly known gas/dust ratio, {\\em (ii)} the sensitivity of \\mdot\\ to the grain size distribution function (we adopt grain sizes from 0.008 to 0.1 micron for the amorphous carbon in the inner shell and constant values in the outer shell), {\\em (iii)} the sensitivity of the emission of the O-rich material to the internal nebular extinction from the C-rich dust (and the ionized gas). The uncertainty in mass loss also reflects on the accuracy of the derived masses, which are $\\sim$ 0.01 and $\\sim$ 0.37 M$_{\\odot}$ for the inner and outer shell respectively. The O-rich shell mass is even more uncertain, as it depends also on the poorly constrained outer boundary radius. The inner radius of the O-rich shell corresponds to a dynamical age of 1050 yr, assuming an expansion velocity of 22 km/sec (Acker et al. 1992). Taking the O-rich mass at face value, this would imply that this shell has been ejected over a time period of approximately 5000 yr. This is considerably shorter than the typical inter-puls period. The mass loss in the O-rich shell is characteristic for OH/IR stars. It may be usefull to search for OH maser emission, as one expects that OH emission remains significant until $\\sim$ 1000 to 1500 yr after the end of the AGB. \\begin{figure} \\psfig{figure=fig3.ps,width=9.0cm} \\caption[model]{Model fit to the ISO-SWS spectrum of BD+30~3639. The long dashed line is the C-rich dust shell, the short-dashed line represents the O-rich detached shell; The dash-dotted line is the sum of both components} \\end{figure} \\subsection{The evolution of [WC] stars} The SWS spectra shown here demonstrate that both BD+30~3639 and He~2-113 were O-rich very recently. The age of the O-rich AGB remnant ($\\sim$~1050~yrs in BD+30~3639, and possibly less for He~2-113) and the mass of the C-rich envelope in BD+30~3639, suggest that the change to C-rich chemistry was triggered by a PAGB pulse and not by a very late thermal pulse. The mass contained in the C-rich dust shell of BD+30~3639 is of the same order of magnitude as the sum of the H and He layers in a young post-AGB star (Bl\\\"ocker et al. 1997). This suggests that the PAGB pulse triggered the removal of the entire outer layer of the star, exposing the He- and C-rich inner layers. Evolutionary calculations of Bl\\\"ocker \\& Sch\\\"onberner (1997) show that in a PAGB thermal pulse, no efficient mixing of the H-rich envelope into the He-burning layers is expected, and hence no [WC] star would form. Only in the case of very late thermal pulses, when the central star is already on the cooling track, mixing of H into the He burning layers efficiently removes H, due to the strongly reduced envelope mass. Since the observations are not consistent with a very late thermal pulse for these two nebulae, we suggest that a PAGB thermal pulse is capable of mixing H and He-rich layers, i.e. contrary to the findings of Bl\\\"ocker \\& Sch\\\"onberner (1997). Recent calculations of Herwig et al. (1997) suggest that a better treatment of convective mixing may result in intershell abundance patterns consistent with the surface C, O, and He abundances of [WC] central stars (see also Leuenhagen \\& Hamann 1997). The latter authors show that late type [WC] stars have small amounts of H and N, which may result from substantial mixing of H and He-rich layers. It is likely that the thermal pulse which caused the change in chemistry is responsible for this mixing. The production of [WC] central stars with O-rich outer shells obviously requires fine-tuning of the timing of the last thermal pulse, either at the very end of the AGB or shortly thereafter. Zijlstra et al. (1991) discovered OH maser emission from IRAS07027-7934, demonstrating that the cool dust is O-rich. Barlow (1998) shows SWS and LWS spectra of CPD-56~8032, also with O-rich dust. This implies that O-rich cool dust is not unusual in nebulae with [WC] central stars. We note that there is evidence that very late thermal pulses can also produce nebulae with [WC] central stars. Pollacco \\& Hill (1994) claimed that the [WC11] star 17514$-$1555 (PN012.2+04.9) is surrounded by both a very compact nebula and an extended, low density nebula, consistent with a very late thermal pulse. This is one of the few late-type [WC] stars which is relatively faint in the IRAS bands. It is unclear at present which mechanism (PAGB thermal pulse with change in chemistry or very late thermal pulse) dominates the formation rate of PNe with [WC] central stars. \\vspace{0.1cm} \\noindent{\\bf Acknowledgements.} LBFMW and AdK gratefully acknowledges financial support from an NWO 'Pionier' grant. FJM acknowledges support from NWO grant 781-71-052. We thank X. Tielens for discussions on PAHs, and H. Dorschner and C. Koike for providing laboratory data of dust species. We thank one of the referees, P. Cox, for communicating results on BD+30~3639 before publication." + }, + "9802/gr-qc9802031_arXiv.txt": { + "abstract": "We study the role played by the compactness and the degree of connectedness in the time evolution of the energy of a radiating system in the Friedmann-Robertson-Walker (FRW) space-times whose $t=const\\,$ spacelike sections are the Euclidean 3-manifold ${\\cal R}^3$ and six topologically non-equivalent flat orientable compact multiply connected Riemannian 3-manifolds. An exponential damping of the energy $E(t)$ is present in the ${\\cal R}^3$ case, whereas for the six compact flat 3-spaces it is found basically the same pattern for the evolution of the energy, namely relative minima and maxima occurring at dif\\/ferent times (depending on the degree of connectedness) followed by a growth of $E(t)$. Likely reasons for this divergent behavior of $E(t)$ in these compact flat 3-manifolds are discussed and further developments are indicated. A misinterpretation of Wolf's results regarding one of the six orientable compact flat 3-manifolds is also indicated and rectified. ", + "introduction": "\\label{Intro} \\setcounter{equation}{0} As general relativity is a purely metrical (local) theory it clearly leaves unsettled the global structure (topology) of space-time. However, in cosmology perhaps the most important problems are related to the global structure of space-time, where the topological degrees of freedom ought to play an essential role. Geometry constrains, but does not determine the topology of a space-time. Consider, for example, the Friedmann-Robertson-Walker (FRW) space-times, whose line element can be given by \\begin{equation} ds^{2} = dt^2 - A^2(t) \\, \\left[\\, \\frac{dr^2}{1 - \\kappa\\, r^2} + r^2\\,(\\,d\\theta^2\\, + \\, \\sin^2 \\theta \\, d\\varphi^{2}\\,)\\, \\right]\\,, \\label{frwds2} \\end{equation} where $A(t)$ is the scale factor, $t$ is the cosmic time, and the constant spatial curvature $\\kappa = 0, \\pm 1$ specifies the type of geometry (flat, elliptic or hyperbolic) of the $t=const$ spacelike section ${\\cal M}_3$. Clearly FRW space-time manifolds ${\\cal M}_4$ can be splitted into ${\\cal R} \\times {\\cal M}_3$. The number of three-dimensional spacelike manifolds ${\\cal M}_3$ which can be endowed with the three possible geometries of the $t=const\\,$ 3-spaces of FRW space-time is quite large~\\cite{Wolf1967,Ellis1971}: for $\\kappa = 0$ there are 18 topologically distinct 3-spaces, while for both $\\kappa = \\pm 1$ an infinite number of 3-spaces exist~\\cite{Ellis1971,LachiezeReyLuminet1995}. Even if we restrict ourselves to orientable and compact manifolds we still have an infinite number of 3-spaces ${\\cal M}_3$ for the elliptic and hyperbolic cases, and six families of topologically dif\\/ferent spacelike manifolds for the flat 3-spaces~\\cite{Wolf1967}~--~\\cite{EllisSchreiber1986}. Since physical laws are usually expressed in terms of local dif\\/ferential equations, in order to be confident about the physical results one derives it is often necessary to have some degree of control over the topological structure of the space-time manifold so as to include constraints imposed by the topology~\\cite{OliveiraReboucasTeixeira1994}. One is then confronted with the question of what topologies are physically acceptable for a given space-time geometry. An approach to this problem is to study the possible observational (or physical) consequences of adopting particular topologies for the space-time~\\cite{Gott1980} -- \\cite{Roukema1996} (see also~\\cite{LachiezeReyLuminet1995} and references therein). In this work we study the role played by the compactness and the degree of connectedness in the time evolution of the energy of a radiating system in the flat FRW space-times whose spacelike $t=const\\,$ sections are endowed with seven different topologies, namely the simply connected 3-space ${\\cal R}^3$, and six multiply connected orientable compact 3-manifolds obtained by suitable identifications of opposite faces of cubes (four) and of hexagonal prisms (two) after suitable turns~\\cite{Ellis1971,EllisSchreiber1986,Gomero1997}. The radiating system we shall be concerned with is represented by a point-like harmonic oscillator (energy source) coupled with a relativistic massless scalar field~% \\cite{SchwalbThirring1964}~--~\\cite{Beig1978}. Similar radiating systems have been used to study a wide class of radiation phenomena as, for example, gravitational waves~\\cite{Unruh1983} or the radiated energy of oscillating electromagnetic dipoles~\\cite{Jackson,Marion}. Our radiating system was also used in the study of dynamic and geometric constraints on the radiation in elliptic FRW expanding universes~\\cite{Bernui1991}~--~\\cite{Beig1978}. In the next section the radiating system is described: the action integral is presented, its variation is performed and the corresponding evolution equations are obtained. We also outline there a proposal for solving these equations. In section~\\ref{GreenFuns} the Green functions of the wave operator are derived for each specific space-time manifold with arbitrary scale factor $A(t)$, and combined with the evolution equations to give the corresponding radiation reaction equations. A misinterpretation~\\cite{Ellis1971}~--~\\cite{EllisSchreiber1986} of Wolf's \\cite{Wolf1967} results (Theorem 3.5.5) regarding one of the six orientable compact flat 3-manifolds is also indicated and rectified therein. In section~\\ref{NumAna}, the radiation reaction equations are numerically integrated. Graphs that show how the system energy varies with the time for each topologically dif\\/ferent flat FRW manifold are presented for static and non-static cases. Our main conclusions are discussed in section~\\ref{Conclusions}. We show that when the $t=const$ sections are the simply connected 3-space ${\\cal R}^3$ the radiation damping phenomenon is present, whereas for all compact flat FRW space-time manifolds we have investigated, the energy $E(t)$ exhibits a few relative minima and maxima followed by a growth of the energy with the time. Possible reasons for this divergent behavior of $E(t)$ in these compact flat 3-manifolds are examined. We also discuss the role played by the compactness and by the degree of connectedness in the energy patterns of our system in these space-time manifolds. Further developments are indicated therein. ", + "conclusions": "\\label{Conclusions} \\setcounter{equation}{0} In this work we have studied the role played by the topological compactness and connectedness in the time evolution of the energy of an harmonic oscillator in flat FRW space-time manifolds, whose $t=const\\,$ sections are (i) the orientable simply connected non-compact 3-space ${\\cal R}^3$ and (ii) six possible flat orientable multiply connected compact 3-manifolds given in table \\ref{FlatTops}. For the ${\\cal R}^3$ case we found that the energy function $E(t)$ exhibits an exponential decay with the time --- the radiation damping ($E(t) \\rightarrow 0$ when $t \\rightarrow \\infty$) takes place, as one could have expected in agreement with~% \\cite{SchwalbThirring1964}~--~\\cite{HoenselaersSchmidty}. For the manifolds ${\\cal T}_1$, ${\\cal T}_2$, ${\\cal T}_3$ and ${\\cal T}_4$ as well as for the manifolds ${\\cal H}_1$ and ${\\cal H}_2$ the behavior of the energy with the time exhibits basically the same pattern: relative minima and maxima occur at dif\\/ferent times for distinct ratios $h/a$ (distinct tessellations) depending on the degree of connectedness of each 3-manifold, and are followed by a growth of $E(t)$. This asymptotical divergent behavior of $E(t)$ for these compact manifolds contrasts with the radiation damping of the energy we have found for ${\\cal R}^3$. There is a quite simple heuristic argument which suports our numerical results, though. If one ad hoc assumes an exponential asymptotical behavior for $Q$, i.e. $Q(t) = \\gamma\\,\\, \\mbox{exp}\\,(\\beta \\, t)$ with $\\beta$ and $\\gamma$ real constants, then for the static cases ($A(t)=1$), and for each of the above compact manifolds, in the limit $t \\rightarrow \\infty$ equation~(\\ref{T1radeq}) reduces to \\begin{equation} \\label{Einfty} \\frac{\\beta^2 + 2\\, \\Gamma\\, \\beta + \\Omega^2}{2\\,\\Gamma} = \\sum_{m=1}^\\infty \\; \\frac{c_m}{a_m}\\,\\, \\exp (-\\beta\\,a_m) \\; , \\end{equation} where $c_m$ is the number of images of the point-like source at a distance $a_m$ in the infinite grid picture. To attain our goal, we will show that~(\\ref{Einfty}) has only one real solution for $\\beta$, which is positive. Indeed, let $f(\\beta)$ be the right hand side of~(\\ref{Einfty}), which is a positive monotone decreasing function of $\\beta$, and such that \\begin{equation} \\lim_{\\beta \\to 0} \\: f(\\beta) = \\infty \\qquad \\quad \\mbox{and} \\qquad \\quad \\lim_{\\beta \\to \\infty } \\: f(\\beta) = 0 \\,. \\end{equation} Thus $f(\\beta)$ lies entirely in the first quadrant of the plane and crosses it from the top-left to the bottom-right. Now, since $\\Gamma > 0$ then for a given pair ($\\Gamma, \\Omega$) the left hand side of (\\ref{Einfty}) is a parabola curved upwards with vertex at $\\beta = - \\Gamma$. Therefore, it always intersects the curve for $f(\\beta)$ in just one point, which is in the first quadrant. In other words, there is only one real $\\beta$ solution to equation~(\\ref{Einfty}), which is positive. The unexpected (unphysical ?) growth of the energy with the time for the above compact flat 3-manifolds cases illustrates that non-trivial topologies can induce rather important dynamic changes in the behavior of a physical system. This type of sensitivity has been refered to as {\\em topological fragility\\/} and can occur without violation of any local physical law% ~\\cite{ReboucasTavakolTeixeira1996}. A rigorous {\\em non-numerical\\/} analysis of the reasons for this surprising divergent behavior of $E(t)$ when compact flat FRW space-times are considered has been carried out, and we hope to publish our results shortly elsewhere. We anticipate, however, that the causes for such a behavior lie in the compactness of the manifold in at least one direction, on the one hand, and in the type of coupling between $Q$ and $\\phi$, on the other hand. A possible physical measure of the degree of connectedness in these {\\em compact\\/} 3-manifolds can be made through the study of the number of emitted rays that return to the origin within a given lapse of time. According to this concept of degree of connectedness one learns from figures 2, 3 and 4 that the greater is the degree of connectedness the earlier is the occurrence of the first relative minimum in the energy function. Incidentally, note that the extension of this concept of degree of connectedness to non-compact 3-manifolds implies that ${\\cal R}^3$ has a null degree of connectedness. This, of course, is indicated in figure 4, which shows a net exponential decay of the energy with the time; no relative minima and maxima come about, which means that no ray returns to the origin. A simple inspection of the graphs for $E(t)$ clearly shows that the derivative $\\dot{E} = \\dot{Q}\\,(\\, \\ddot{Q} + \\Omega^2\\, Q\\,)$ of the energy function is discontinuous at a few points. Indeed, for the static case, for example, using (\\ref{T1radeq}) one obtains \\begin{equation} \\dot{E}(t) = 2 \\, \\Gamma \\, \\dot{Q}(t) \\left[ \\, \\sum_{n_{1},n_{2},n_{3}} \\frac{1}{a(\\vec{n})}\\,\\, \\Theta(t-a(\\vec{n})) \\,\\, Q(t - a(\\vec{n})\\,) \\,-\\, \\dot{Q}(t) \\right] \\,. \\end{equation} From this equation one sees that the discontinuities occur at $t=a(\\vec{n})$, that is, they come about each time a new term $ Q(t-a(\\vec{n}))\\,/\\,a(\\vec{n})$ is taken into account in the right hand side of (\\ref{T1radeq}). A question which naturally arises here is whether the inverse problem, i.e. that of determining the basic cell (topology) corresponding to the spacelike $t=const$ sections from the graphs of the energy $E(t)$, can be solved. Regarding this problem it is clear that one can find the distances of the point-like source to its images by using where the discontinuities of $\\dot{E}(t)$ take place, and the number of images at a given distance through the magnitude of the corresponding discontinuities. So, one can probe the topology of the 3-spaces at least in a few cases. It is not yet clear whether an algorithm for solving the inverse problem for the most general (Euclidean) setting can be found, though. As far as we are aware~\\cite{LachiezeReyLuminet1995} this is the first work in which a physical consequence of adopting the flat hexagonal prisms ${\\cal H}_1$ and ${\\cal H}_2$ has been studied. It is of worth to emphasize that when the expansion of the universe is considered the degree of connectedness is less than the ones for the static cases. For the manifold ${\\cal T}_{1}$ this can be seen by comparing figure 2 (static case $\\dot{A}(t) = 0$) and figure 5 (monotone expansions $\\dot{A}(t) > 0$). Before closing this article we should like to stress that our study does not cover all possible spatially compact orientable flat FRW manifolds. Thus, for example, we have not considered that for the 3-manifolds ${\\cal T}_1$ and ${\\cal T}_3$ the basis of the basic cell need not to be a square, it can be a parallelogram. The restriction we have made, however, does not seem to be decisive for the patterns of the behavior of the energy with the time we have found. To conclude we remark that the study of radiation damping in elliptic ($\\kappa = 1$) FRW manifolds in which the $t=const$ sections are endowed with dif\\/ferent (orientable compact) topologies is being carried out. \\vspace{5mm}" + }, + "9802/astro-ph9802076_arXiv.txt": { + "abstract": "Period changes of $\\delta$ Scuti stars have been collected or redetermined from the available observations and are compared with values computed from evolutionary models with and without convective core overshooting. For the radial pulsators of Pop. I, the observations indicate (1/P) dP/dt values around 10$^{-7}$ year$^{-1}$ with equal distribution between period increases and decreases. The evolutionary models, on the other hand, predict that the vast majority should show increasing periods. This increase should be a factor of about ten times smaller than observed. For nonradial $\\delta$ Scuti pulsators of Pop. I, the discrepancies are even larger. The behavior suggests that for these relatively unevolved stars the rate of evolution cannot be deduced from the period changes. The period changes of most Pop. II $\\delta$ Scuti (SX Phe) stars are characterized by sudden jumps of the order of $\\Delta P/P \\sim$ 10$^{-6}$. However, at least one star, BL Cam, shows a large, continuous period increase. The variety of observed behavior also seems to exclude an evolutionary origin of the changes. Model calculations show that the evolutionary period changes of pre-MS $\\delta$ Scuti stars are a factor of 10 to 100 larger than those of MS stars. Detailed studies of selected pre-MS $\\delta$ Scuti stars are suggested. ", + "introduction": "$\\delta$ Scuti variables are short-period pulsators situated inside the classical instability strip on and near the main sequence (hereafter called MS). They represent stars in the hydrogen-burning stage with masses from 1.5 $M_{\\odot}$ to about 2.5 $M_{\\odot}$ for the more evolved variables. The majority of $\\delta$ Scuti stars pulsates with a multitude of nonradial p-- and g--modes (often of mixed character) and low amplitudes near the detection limit of 0.5 millimag. A numerically small subgroup has exceptionally low rotational velocities of $v \\leq $ 30 km s$^{-1}$ with pulsation properties resembling the classical cepheids and RR Lyrae stars: radial pulsation in the fundamental or first overtone mode with high amplitudes between 0.3 and 1 mag. In the past, the stars in this high-amplitude (HADS) subgroup were sometimes referred to as Dwarf Cepheids or AI Vel stars. They are observationally favored for the study of period stability because of large amplitudes and the presence of only a few excited modes (usually one or two). The period changes caused by stellar evolution in and across the Lower Instability Strip permit an observational test of stellar evolution theory, provided that other physical reasons for period changes can be excluded. The period -- mean density relation for pulsating stars predicts a period-luminosity-color-mass relation of the form \\[\\log P = -0.3 M_{\\rm bol} - 3 \\log T_{\\rm eff} - 0.5 \\log M + \\log Q + 12.708\\] where P is the radial period and $M$ is the stellar mass in solar units. For the $\\delta$ Scuti variables with known radial pulsation, the observed coefficients are in almost exact agreement with those predicted by the period--mean density relation (Breger 1980). This confirmation of the theoretical relation is very important: since the period-luminosity-relation holds for stars in different evolutionary stages, it should also hold for a single star at different evolutionary stages. For individual stars, evolutionary period changes must occur, at least over long time scales. An evolutionary change in $T_{\\rm eff}$ and $M_{\\rm bol}$ leads to a period change of size \\[\\frac{1}{P} \\frac{dP}{dt} = -0.69\\frac{dM_{\\rm bol}}{dt} - \\frac{3}{T_{\\rm eff}} \\frac{dT_{\\rm eff}}{dt} + \\frac{1}{Q} \\frac{dQ}{dt},\\] where the coefficient of 0.69 has been derived by multiplying the coefficient of 0.3 by $\\ln$ 10 from the differentiation. In the Lower Instability Strip, where the $\\delta$ Scuti stars are found, stellar evolution leads to increasing periods in the vast majority of stars, with predicted increases of (1/P)dP/dt from 10$^{-10}$ year$^{-1}$ on the MS to 10$^{-7}$ year$^{-1}$ for the longer-period evolved variables. (More exact values will be computed in a later section of this paper.) Such period changes should be observable. There have been several comparisons between theoretical and observed period changes. The careful compilation by Percy et al. (1980) led to inconclusive results. They state that `the study of period changes has had a long and sometimes dubious history' and find that `in most cases is a source of confusion and frustration.' Breger (1990a) had more data at his disposal and noted that for the four evolved $\\delta$ Scuti stars of Pop. I the observed period decreases were in contradiction to the radius increases predicted from stellar evolution. Guzik \\& Cox (1991) have examined possible explanations for the period decreases seen in evolved $\\delta$ Scuti stars. While their models indeed predicted smaller period increases, period decreases were not found. The fact that all four evolved $\\delta$ Scuti stars show period decreases may conceivably be a reflection of an additional period changing mechanism (leading to random changes) coupled with small-number statistics. The recent results for the star AI Vel (Walraven et al. 1992) support such a view. They have found that only one of the two radial frequencies shows a period increase, while the second frequency was essentially constant. The example of AI Vel shows that (slow) stellar evolution is not the only mechanism to generate changes in the periods of pulsation. This means that for an individual star the conversion of observed period changes into stellar evolution rates (e.g. radius changes) has to be applied with caution. The nonevolutionary period changes should cancel out only for a larger group of stars in the same stage of evolution. An important update and analysis was provided by Rodriguez et al. (1995), who confirmed the observed period decreases for evolved $\\delta$ Scuti stars. They also found relatively large period increases in excess of the expected values for the less evolved Pop. I $\\delta$ Scuti variables and decreasing periods for Pop. II variables. Furthermore, their study provided further evidence for the earlier suspicions that the observed period changes are not caused by evolution alone. The present paper examines period changes in $\\delta$ Scuti stars from a theoretical point of view, provides a review of the observational status, and compares the two. ", + "conclusions": "" + }, + "9802/astro-ph9802130_arXiv.txt": { + "abstract": "The fate of the cooling gas in the central regions of rich clusters of galaxies is not well understood. In one plausible scenario clouds of atomic or molecular gas are formed. However the mass of the cold gas, inferred from measurements of low energy X-ray absorption, is hardly consistent with the absence of powerful CO or 21 cm lines emission from the cooling flow region. Among the factors which may affect the detectability of the cold clouds are their optical depth, shape and covering fraction. Thus, alternative methods to determine the mass in cold clouds, which are less sensitive to these parameters are important. For the inner region of the cooling flow (e.g. within the radius of $\\sim$50--100 kpc) the Thomson optical depth of the hot gas in a massive cooling flow can be as large as $\\sim 0.01$. Assuming that the cooling time in the inner region is few times shorter than the life time of the cluster, the Thomson depth of the accumulated cold gas can be higher accordingly (if most of the gas remains in the form of clouds). The illumination of the cold clouds by the X-ray emission of the hot gas should lead to the appearance of a 6.4 keV iron fluorescent line, with an equivalent width proportional to $\\tau_T$. The equivalent width only weakly depends on the detailed properties of the clouds, e.g. on the column density of individual clouds, as long as the column density is less than few $10^{23}~cm^{-2}$. Another effect also associated exclusively with the cold gas is a flux in the Compton shoulder of bright X-ray emission lines. It also scales linearly with the Thomson optical depth of the cold gas. With the new generation of X--ray telescopes, combining large effective area and high spectral resolution, the mass of the cold gas in cooling flows (and it's distribution) can be measured. ", + "introduction": "The radiative cooling time of the gas in the central parts of rich clusters of galaxies can be considerably shorter than the Hubble time \\cite{lea73,sil76,cb77,fn77} and the gas may cool down below X--ray temperatures, forming a cooling flow (see Fabian, 1994 for review). Strong peaks in the surface brightness distribution, observed in many clusters of galaxies (e.g. Branduardi--Raymont et al. 1981, Fabian et al. 1981, Canizares, Stewart \\& Fabian A.C. 1983, Fabricant \\& Gorenstein 1983, Jones \\& Forman 1984, Stewart et al. 1984), are usually considered as indicators of a cooling flow. Evidence for cool gas was also found in the spectroscopic observations, which revealed the presence of emission lines, characteristic of gas with the $T_e \\le 5\\times 10^6$ K gas (e.g. Canizares et al., 1979). Detailed studies of the surface brightness distribution in the cooling flows lead to the conclusion, that a fraction of the gas drops out from the flow at different radii, such that the mass deposition rate is proportional to the radius: $\\dot{M} (1$ or the steady-state model with $q_0 = -1$, can be addressed with existing number count data. Although also discussed by Sandage \\markcite{sandage61ev}(1961b), the effects of luminosity evolution of galaxies on the $n(m)$ relation, were not studied in detail until the 1970s. In a series of papers culminating in a review of her galaxy count model (Tinsley \\markcite{tinsley77}1977; and references therein) developed the technique of population synthesis, in which she took an initial mass function (IMF) of stars, and observed broadband stellar colors to predict the color and luminosity evolution of stellar populations. To construct her model, she made several assumptions which would dominate the study of galaxy counts for the next two decades. She assumed that all galaxies formed at the same time, and that the time-scale of star-formation was the factor which determined the relation between morphological Hubble type and SED. In this picture, elliptical galaxies formed their stars in a short initial burst, while spiral galaxies had longer timescales of star-formation. She assumed that there was conservation of galaxy number, and merging was not an important factor. The population synthesis approach dominated studies of galaxy evolution in the 1980s, when workers used stellar spectra instead of broadband colors to construct galaxy templates. In a series of papers Bruzual \\markcite{bruzual83}(1983a) extended the technique to the ultraviolet (UV) using new observations of stars with the IUE telescope. At about the same time, Arimoto \\& Yoshii \\markcite{arimoto86}(1986), 1987 and Guiderdoni \\& Rocca-Volmerange \\markcite{guiderdoni87}(1987) included the effects of chemical evolution of galaxies, tracing the change in metallicity as the population of stars within the galaxy evolves. These results were included in galaxy count models by Yoshii \\& Takahara \\markcite{yt88}(1988) and Guiderdoni \\& Rocca-Volmerange \\markcite{guiderdoni90}(1990). \\subsection{Galaxy Count Models} The galaxy count model of Yoshii \\& Takahara \\markcite{yt88}(1988) serves as the basis for this paper, because many of the ingredients for the model have now become standard. They began with a measured local luminosity function. They assumed that the characteristic magnitude ($M^*$ in the Schechter \\markcite{schechter}1976 parameterization) and faint end slope ($\\alpha$) of the luminosity function was independent of Hubble type in the $b_J$ filter, and set the distribution of types (and thus $\\phi^*$ as a function of type) to match those observed in a local survey. They then converted that luminosity function to other filters using type-dependent local galaxy colors from their population synthesis model. This has the effect that while $\\alpha$ is the same for each type of galaxy in each filter, $M^*$ is a function of type in filters other than $b_J$, and the summed luminosity function in other filters has an overall $\\alpha$ that is different from that of the individual types. Examples of this effect are given below. Other galaxy count modellers (Guiderdoni \\& Rocca-Volmerange \\markcite{guiderdoni90}1990) used type-dependent luminosity functions (e.g. King \\& Ellis \\markcite{kingellis}1985), even though these were of necessity less accurately determined. The population synthesis models give the SED of each type of galaxy as a function of time since the galaxy formed. (Following Tinsley \\markcite{tinsley77}1977, the morphological galaxy types are each equated with a star--formation history). The chosen cosmology give the relation between redshift, distance modulus and K--corrections, and the time-redshift relation is applied to the luminosity evolution model so as to include evolutionary corrections. The model is then integrated over redshift to give the $n(m)$ relation, or integrated over other parameters to give $n(z)$, and color distributions. Recently a new type of model has appeared. The semi-analytic approach, motivated by cosmological theory, was proposed by White \\& Frenk \\markcite{white91}(1991) and Cole \\markcite{cole91}(1991), and developed further by Lacey et al.\\ \\markcite{lacey93}(1993), Kauffman, White \\& Guiderdoni \\markcite{kauffmann93}(1993), Kauffman, Guiderdoni \\& White \\markcite{kauffmann94}(1994) and Cole et al.\\ \\markcite{cole}(1994). This type of theory begins with the physics of the big bang, cold dark matter, and hierarchical structure formation, and traces the formation of galaxies {\\em ab initio} using a semi-analytic approach constrained by numerical simulations. It then predict observables, such as the number counts, luminosity functions, colors, redshift distributions, and the Tully--Fisher relation. Baugh, Cole \\& Frenk \\markcite{baughmorph}(1996) have extended this model to study the distribution of morphological Hubble types, as measured by the bulge-to-disk ratio. The fundamental difference between this type of model and the traditional number count models is the direction in time. Traditional models take the local population and attempt to extend it backwards in time, and determine what it evolved from. The semi--analytic approach takes the initial conditions set by the big bang and attempts to trace their evolution into the local population. While the semi--analytic theory is the most physically motivated, and may eventually provide the greatest understanding of galaxy formation and evolution, there is still a place for the traditional models. Traditional models are still the best way of comparing different data sets, and this is the approach I take in this paper. \\subsection{Galaxy Count Data} A considerable amount of observational galaxy count work was done in the 1980s using photographic plates, both on Schmidt telescopes to get bright counts, and on $4m$ class telescopes to get faint counts. The modern era of galaxy count data, using linear detectors (CCDs) to count galaxies to very faint levels began with the work of Tyson \\markcite{tyson88}(1988) and Lilly, Cowie \\& Gardner \\markcite{lcg}(1991). Those workers established the counts in several optical filters to $B\\simeq28$, while a series of papers by Metcalfe et al.\\ \\markcite{metcalfe91}(1991) established CCD counts at intermediate magnitudes. These data, when compared to the modeling work described above, led to what has become known as the faint blue galaxy problem. The data showed an excess in the number of galaxies over the no-evolution model predictions. While the initial interpretation was that the surveys were detecting the ultraviolet light from star-forming galaxies in an open cosmology, further investigation, including redshift surveys (Broadhurst, Ellis \\& Shanks \\markcite{bes}1988; Colless et al.\\ \\markcite{colless}1990; Cowie, Songaila \\& Hu \\markcite{csh91}1991) showed that this excess was not at high redshift. The excess was not seen in the near-infrared $K-$band number counts (Gardner, Cowie \\& Wainscoat \\markcite{gcw}1993), and was made up of blue galaxies. The spatial resolution of images taken by the Hubble Space Telescope (HST) has allowed the morphological study of faint galaxies. The HST Medium Deep Survey key project (Griffiths et al.\\ \\markcite{griffiths}1994; hereafter MDS) has proven to be a great success in characterizing the optical morphology of galaxies at intermediate magnitudes. In conjunction with the fainter images of the Hubble Deep Field (Williams et al.\\ \\markcite{williams}1996; hereafter HDF), the MDS has revealed an excess of galaxies with unusual or disturbed morphologies, over what is predicted by the number count models (Glazebrook et al.\\ \\markcite{kgbmds}1995; Driver, Windhorst \\& Griffiths \\markcite{drivermds}1995; Abraham et al.\\ \\markcite{abraham}1996). These irregular galaxies likely make up the population of excess faint blue galaxies. Bright $b_J-$band galaxy counts, measured on scanned photographic plates covering 4300 square degrees, show a steeper slope than is predicted by the models (Maddox et al.\\ \\markcite{maddox}1991). This steep slope was initially believed to be due to rapid evolution at low redshift, but recently the data has been called into question. Systematics in the photometry, either due to the rapid plate scanning process, (Metcalfe, Fong \\& Shanks \\markcite{metcalfeapm}1995), or to the distribution of central surface brightness in the galaxies, (McGaugh \\markcite{mcgaugh}1994), could steepen the slope of the number counts artificially. This has caused considerable debate about the magnitude of the faint blue galaxy problem. When the models are normalized at the bright end of the local steep slope, $B\\simeq 15$, the excess in the counts is a factor of $2$ to $4$ over the model predictions at $B\\simeq22$. The results of redshift surveys of galaxies selected at this magnitude level, however, show little departure from no-evolution predictions of the shape of the number-redshift relation (Colless et al.\\ \\markcite{colless}1990), although their absolute numbers, or normalization showed the excess seen in the number counts. If the models are normalized at the faint end of the local steep slope, $B\\simeq18$, and passive evolution is included in the model predictions, then the faint blue galaxy excess does not appear until $B>25$ for the closed universe case, which is fainter than any of the spectroscopic redshift surveys. Recent bright number counts, obtained with CCD imaging of 10 square degrees, do not confirm the steep slope, and are consistent with the higher normalization (Gardner et al.\\ \\markcite{gardnerklfnc}1996). Larger CCD surveys are needed to confirm this result. There have been several attempts to modify the traditional number count models to explain the faint blue galaxy problem by introducing one or more additional free parameters into the model. Rocca-Volmerange \\& Guiderdoni \\markcite{rocca91}(1990) and Broadhurst, Ellis \\& Glazebrook \\markcite{beg}(1992) relaxed the constraint of conservation of galaxy number, and introduced large amounts of merging. A disappearing dwarf galaxy model (Dekel \\& Silk \\markcite{dekel}1986; Cowie et al.\\ \\markcite{csh91}1991; Babul \\& Rees \\markcite{babul}1992) has gained support from deep redshift surveys (Lilly et al.\\ \\markcite{cfrs}1995; Ellis et al.\\ \\markcite{autofib}1996) which see a steepening of the faint end slope of the luminosity function to higher redshift, particularly for the bluer galaxies. High normalization models have much less of a faint blue galaxy excess to explain, and succeed in fitting most of the observations with modifications to the luminosity evolution models. One recent attempt used a steep slope for the luminosity function of late-type spiral galaxies, which has an effect on the counts similar to the disappearing dwarf models (Metcalfe et al.\\ \\markcite{metcalfenat}1996). Cowie et al.\\ \\markcite{cowie96}(1996) recently proposed a downsizing model, in which less luminous galaxies evolve rapidly at low to intermediate redshift, while bright galaxies evolve rapidly at redshift $z>1$. These models have usually attempted to fit all of the available field galaxy survey data, and do so by introducing additional free parameters into the models. Fitting the data, however, is often only done to within a factor of $\\sim 2$. In this paper I take a more modest approach. Rather than trying to construct a single model which fits all of the data, I construct an extendable or general purpose model. The model can be tuned to fit one observation, and then used to predict an observation in another filter. This is essentially a translation program, translating observations between filters. Number counts and luminosity functions measured in different surveys can be directly compared. When the filters are similar, such as the ground-based $U$ band filter and the WFPC2 $U_{300}$ filter, this can be done in a relatively model independent way, as a model which fits the $U-$band number counts is likely to also fit the $U_{300}$ counts. When the filters have very different central wavelengths, such as $B$ and $K$, or when the observations are at very different magnitudes, then the comparison will depend more strongly on the model parameters chosen. I also will use this program to predict the results of future ground and space mission surveys, by extrapolating the most relevant existing data into the filters and depths that will be achieved. In section 2 I will outline the equations and parameters that make up the model. In section 3 I will use the model to compare existing survey data, and to predict the results of future surveys. Section 4 is a summary. ", + "conclusions": "Galaxy counts are one of the classical cosmological tests, but their interpretation remain difficult. Model predictions of the counts, and the color and redshift distributions of galaxies are subject to uncertainties in the spectral energy distributions and evolution of galaxies, and in free parameters specifying the luminosity functions, cosmological geometry, the number and distribution of galaxy types, and the effects of dust and merging. However, there is a growing body of observational data that constrains these models. In this paper, I have presented a model in which the free parameters can be easily adjusted, and which can be extended to include new parameters. The greatest power of this model is in converting observational data taken in one filter into another filter to compare with other data, and extrapolating to longer or shorter wavelengths. The software, parameter files and data presented in this paper are available at http://hires.gsfc.nasa.gov/$\\sim$gardner/ncmod. \\bigskip I wish to acknowledge funding by the Space Telescope Imaging Spectrograph Investigation Definition Team through the National Optical Astronomical Observatories, and by the Goddard Space Flight Center. \\clearpage" + }, + "9802/astro-ph9802166_arXiv.txt": { + "abstract": "We present a two dimensional map of the gas temperature distribution in the Centaurus cluster, based on ASCA observations derived using a novel approach to account for the energy dependent point spread function. Along with a cool region, centered on NGC4696, asymmetric temperature variations of moderate amplitude are detected. The hottest region roughly coincides with the position of the second brightest galaxy NGC4709, known to be the dominant galaxy of one of the subgroups in the Centaurus cluster. ROSAT images show faint surface brightness emission also centered on this galaxy. The imaging and spectral results suggest that a subcluster, centered on NGC4709 is merging with the main cluster centered on NGC4696, in agreement with the earlier suggestion by \\cite{lcd86} that these two systems are located at the same distance despite their different line-of-sight velocities. ", + "introduction": "The Centaurus cluster is a nearby (mean velocity of the dominant component is about 3000 \\kms, e.g. Lucey, Currie, \\& Dickens 1986) X-ray bright cluster with a 2--10 keV luminosity of $7\\times 10^{43}$~ergs~s$^{-1}$ (Jones \\& Forman 1978). The cluster X--ray emission is strongly peaked on the cD galaxy NGC4696, indicating a cooling flow with a moderate mass deposition rate of $\\sim$ 30--50 $M_\\odot$ (e.g. Matilsky, Jones \\& Forman 1985; Allen \\& Fabian 1994; White, Jones \\& Forman 1997). The cluster was intensively studied in X-rays with imaging instruments by Einstein (Matilsky, Jones \\& Forman 1985), ROSAT (Allen \\& Fabian 1994), and ASCA (Fukazawa et al.\\ 1994). Rather smooth X-ray isophotes, having a slightly elliptical shape (Allen \\& Fabian 1994), and the presence of a cooling flow suggest that the Centaurus cluster is a relatively relaxed system. However, measurements of line-of-sight velocities indicate the presence of two distinct components (Lucey et al. 1986), one centered on NGC4696 and another dominated by NGC4709, but having a factor of 2.5 fewer galaxies. These two components, denoted as Cen 30 and Cen 45, have mean velocities and dispersions of ${\\rm v}=3041$\\kms, $\\sigma=586$\\kms~ and ${\\rm v}=4570$\\kms, $\\sigma= 280$\\kms~ respectively (Lucey et al. 1986). Although this may be the simple projection of two spatially separated groups, \\cite{lcd86} argued, on the basis of comparison of luminosity functions, color--magnitude relations and galaxian-radius distributions, that both components are in fact located at the same distance. This would imply that the main cluster, Cen 30, is in the process of accreting the smaller group, Cen 45, at the present time. We show below that the ASCA observations provide strong support that the Centaurus cluster is undergoing a merger. The structure of the paper is as follows: Section 2 briefly describes the X--ray observations and data reduction procedure. In Section 3 we discuss results of the analysis and present arguments in favor of an ongoing merger of the Cen 30 and Cen 45 subclusters. Section 4 summarizes the results. ", + "conclusions": "Despite of the large difference ($\\sim 1500$\\kms) in the velocities of the Cen 30 and Cen 45 groups, \\cite{lcd86} presented strong arguments that these two groups are located at the same distance. Luminosity functions, color--magnitude relations and galaxian-radius distributions support the assumption of a common distance and disagree with the contention that Cen 45 is significantly more distant than Cen 30, in accordance with their line-of-sight velocities. Dressler (1995) came to the same conclusion studying surface brightness fluctuations of several galaxies from both groups. These results imply that the smaller Cen 45 group is merging with the dominant Cen 30 group. The analysis of X--ray data further supports this scenario. Indeed, if one were observing the projection of two independent clusters, we would expect the overlapping region to have a lower temperature due to the predicted lower temperature of the smaller group Cen 45. Instead, we find that the temperature is highest in the region close to the dominant galaxy of the Cen 45 group, which we interpret as heating due to the interaction of the two subgroups. Note however that the presence of the cooling flow around NGC4696 restricts our ability to correctly estimate the extent of the hot region. We might also expect that when a poor group is entering a dense cluster environment for the first time, the relatively cold gas will be stripped due to the high ram pressure. The marginally significant filament near NGC4696B, seen in the ROSAT PSPC images, might be explained as the result of such an interaction. Recently \\cite{sjf97} found, in the $0.75$\\deg circle around NGC4696, a very large velocity dispersion for Cen 30 ($933\\pm 118$\\kms) and a very low dispersion for Cen 45 ($131\\pm 43$\\kms). The assumption that the main cluster has a velocity dispersion of 933\\kms~ would make the so--called $\\beta$-problem very extreme. This parameter (which characterizes the ratio of energy per unit mass for galaxies to that in the gas) can be estimated from the X--ray surface brightness distribution, $\\beta_f$, or calculated from the observed velocity dispersion and gas temperature, $\\beta_s$. The slope of the PSPC surface brightness profile implies values of $\\beta_f\\sim$ 0.4--0.45. However, $\\beta_s=\\frac{\\mu m_p \\sigma_r^2}{kT}\\sim$1.4 for $T\\sim$ 4 keV and $\\sigma_r=$933\\kms. Even though these values for $\\beta$~ disagree for many clusters, for Centaurus this disagreement is far larger than usual. An alternative and more likely explanation for the large velocity dispersion is that a large scale filament, which defines the preferential direction of accretion onto the cluster, lies close to the line of sight. This would explain the high observed velocity difference of the infalling group and would also account for the large velocity dispersion observed by \\cite{sjf97}. The presence of a large scale filament crossing the center of Cen~30 also would naturally explain the alignments already mentioned above. Thus, accretion along a filament, lying mostly along the line of sight, but also partially along the direction defined by the centers of Cen~30 and Cen~45, would explain the common alignments of 1) the major axis of NGC4696, 2) the filament emanating from NGC4696B and 3) the nearby chain of groups (Lucey et al. 1986). The effects of large scale filaments are seen around other clusters. For example, observations of A1689 suggest the presence of a large scale filament directly along the line of sight (Daines et al. 1998). In A1689, the three brightest galaxies in the direction of the cluster lie within $20''$ of the cluster center, only one has a velocity equal to the cluster mean while the other two have velocity differences of +4767 \\kms and -2686 \\kms, far too large to be explained by the velocity dispersion of a relaxed $\\sim10$ keV temperature cluster. The nearby cluster A85 also shows the common alignment of several galaxy concentrations/subclusters with the major axis of the central cD galaxy (Durret et al. 1997). The X-ray temperature map for A85 (Markevitch et al. 1998) shows a hot region, suggestive of a shock produced by one of the infalling subclusters. Large scale filaments and their effects on clusters also are found in recent numerical simulations which have adequate resolution. Colberg et al. (1998) showed that clusters accrete matter from a few preferred directions, defined by filamentary structures, and that the accretion persists over cosmologically long times. Thus, it should not be surprising to see common alignments between central galaxy major axes, infalling subclusters, surrounding groups, and, even on the scale of neighboring clusters (West et al. 1995). With the availability of cluster temperature maps, the X-ray observations allow us to differentiate between superpositions (chance alignments) and true mergers where shock heating occurs. Hence, as shown in this paper for Centaurus, we can now study the details of the merger process which provide strong evidence for accretion along preferred directions as predicted by numerical simulations of hierarchical structure formation." + }, + "9802/astro-ph9802093_arXiv.txt": { + "abstract": "We live in a dusty Universe, and correcting for the dust extinction and reddening affects almost all aspects of the optical astronomy. Recently Schlegel, Finkbeiner \\& Davis published an all-sky reddening map based on the COBE/DIRBE and IRAS/ISSA infrared sky surveys. Their map is intended to supersede the older Burstein \\& Heiles reddening estimates. In this paper I test this new reddening map by comparing the reddening values for a sample of 110 $|b|>5\\deg$ Galactic globular clusters selected from compilation of Harris. I find a good agreement for globular clusters with galactic latitude $|b|>20\\deg$ and fair overall agreement for globular clusters with $20>|b|>5\\deg$, but with several significant deviations. I discuss four individual clusters with largest deviations, NGC 6144, Terzan 3, NGC 6355 and IC 1276, in order to investigate the reasons for these large deviations. It seems that the new reddening map overestimates the reddening in some large extinction regions. However, with its high spatial resolution the new reddening map can be used to estimate the relative variation of the reddening on scales $\\lesssim 10'$. ", + "introduction": "We live in a dusty Universe (Hoover 1998, private communication), and correcting for the dust extinction and reddening affects almost all aspects of optical astronomy. For us, observing from within the Milky Way, it is of crucial importance to know how much Galactic dust there is towards various objects. Burstein \\& Hailes (1982; hereafter: BH) constructed an all-sky reddening map, used extensively by the astronomical community.\\footnote{Their paper was cited 540 times between 1992 and 1997} Recently, Schlegel, Finkbeiner \\& Davis (1998; hereafter: SFD) published a new all-sky reddening map, based on the COBE/DIRBE and IRAS/ISSA maps.\\footnote{The reddening map and related files and programs are available using the {\\tt WWW} at: {\\tt http://astro.berkeley.edu/davis/dust/}} This map is intended to supersede the BH map in both the accuracy (16\\%) and the spatial resolution ($6.1'$). Indeed, the potential of the SFD reddening map is immediately apparent after examining their Fig.7. It is therefore important to independently test the SFD map to determine the accuracy of the predicted reddenings. In this paper I use an electronic catalog of the Galactic globular clusters compiled by Harris (1996) to test the SFD map. In Section~2 I compare the values of reddening and discuss the dependence of the deviations on various properties of the GCs. In Section~3 I discuss several individual GCs for which the derived deviations are the largest. ", + "conclusions": "\\begin{figure}[t] \\plotfiddle{fig3.ps}{9.6cm}{0}{60}{60}{-190}{-105} \\caption{Deviations between the values of the reddening $E(B-V)_H$ obtained using the electronic catalog of Harris (1996) and the reddening $E(B-V)_{BH}$ obtained using the reddening map of Burstein \\& Hailes~(1982). In the upper panel I plot the deviations $\\Delta E(B-V)\\equiv E(B-V)_{BH}-E(B-V)_H$ as the function of $E(B-V)_H$. Larger symbols denote clusters with $|b|>20\\deg$. In the lower panel I plot the histogram of the $\\Delta E(B-V)$ values for 82 Galactic globular clusters with $|b|>10\\deg$ and for the subset of 46 GCs with $|b|>20\\deg$ (shaded).} \\label{fig:bh} \\end{figure} It might be also interesting to compare the deviations between the values of $E(B-V)_H$ and those predicted by the BH map, which I denote $E(B-V)_{BH}$. In Figure~\\ref{fig:bh}, analogous to Figure~\\ref{fig:sfd}, I plot the deviation $\\Delta E(B-V)\\equiv E(B-V)_{BH}-E(B-V)_H$. Since the BH map excludes objects with $|b|<10\\deg$, there is fewer (82) GCs now, with smaller range of $E(B-V)_H$. There is a clear trend between $E(B-V)_H$ and $\\Delta E(B-V)$, indicating that the BH map underestimates the value of reddening, especially in high reddening regions. As noticed by SFD, even in high galactic latitude regions there is a systematic offset of $\\sim 0.02\\;$mag between the SFD and the BH maps, with the BH map predicting lower values of the reddening. Comparing Figures~\\ref{fig:sfd} and \\ref{fig:bh}, it is clear that overall the SFD map does a much better job in predicting the reddening values of the Galactic globular clusters. To summarize, the SFD map predicts well the reddening for most of the globular clusters in the Harris (1996) sample, especially for the $|b|>20\\deg$ clusters. It is however possible that it overestimates the reddening in some large extinction regions, such as in the case of NGC 6144, NGC 6355 and IC 1276 discussed above, although it might be that the $E(B-V)$ estimates based on the integrated properties of globular clusters are at fault here (see discussion in Burstein \\& Hailes 1978). In any case, the SFD map is a very valuable tool in predicting large reddening gradients on scales of $\\lesssim 10'$. Further tests of the SFD map would be most useful, using for example field RR Lyrae stars (Burstein \\& Hailes 1978), such as the sample of Layden (1998), but with reddenings not based on the BH map." + }, + "9802/astro-ph9802087_arXiv.txt": { + "abstract": "The evolution of neutron stars in close binary systems with a low-mass companion is considered assuming the magnetic field to be confined within the solid crust. We adopt the standard scenario of the evolution in a close binary system in accordance with which the neutron star passes throughout four evolutionary phases (\"isolated pulsar\" -- \"propeller\" -- accretion from the wind of a companion -- accretion due to Roche-lobe overflow). Calculations have been performed for a great variety of parameters characterizing the properties both of the neutron star and low-mass companion. We find that neutron stars with more or less standard magnetic field and spin period being processed in low-mass binaries can evolve to low-field rapidly rotating pulsars. Even if the main-sequence life of a companion is as long as $10^{10}$ yr, the neutron star can maintain a relatively strong magnetic field to the end of the accretion phase. The considered model can well account for the origin of millisecond pulsars. {\\bf Key words}: pulsars: recycled - neutron stars: structure - neutron stars: magnetic fields - X-ray binaries ", + "introduction": "In recent years evidences have been obtained for a relatively fast magnetic field decay in neutron stars undergoing accretion in binary systems (see, e.g., Taam \\& van den Heuvel 1986, Bhattacharya \\& van den Heuvel 1991). The weak surface magnetic fields of many binary radiopulsars ($B \\leq 10^{9}$G), for which an efficient accretion has taken place, strongly suggest the idea of the accretion induced magnetic field decay in these neutron stars. Most of the low-magnetic-field pulsars are believed to be old neutron stars experienced a mass transfer in low-mass binaries where accretion ensues due to Roche-lobe overflow and may last relatively long. Recently Geppert \\& Urpin (1994) and Urpin \\& Geppert (1995) have suggested a simple and efficient mechanism of a rapid field decay in neutron stars undergoing accretion. Accretion changes drastically the thermal evolution of the neutron star if the accretion rate is sufficiently high and a mass transfer lasts a sufficiently long time (see, e.g., Fujimoto et al. 1984). The energy release due to accretion heats the neutron star and reduces the crustal conductivity since the latter depends on the temperature. A decrease of the conductivity accelerates the ohmic decay of the magnetic field if this field is maintained by currents in the crust. If the mass transfer phase lasts as long as $10^{6}-10^{7}$ yr the surface field strength can be reduced by a factor $\\sim 10^{2}-10^{4}$ depending on the magnetic configuration at the beginning of the accretion phase. Accretion influences not only the field decay but the spin evolution of pulsars as well, and the strength of the magnetic field is one of the dominating factors that determines this evolution. If the total amount of the accreted matter is large enough ($\\geq 0.01 - 0.1 M_{\\odot}$), the neutron star in a low-mass binary can be \"recycled\" to millisecond periods (Alpar et al. 1982). The accretion phase, however, does not exhaust all possible stages in the evolution of neutron stars entering low-mass binary systems. The evolution of such systems is extremely long since even the main-sequence lifetime of a low-mass companion exceeds $10^{9}$ yr. In accordance with the standard scenario, the neutron star in a close binary can pass throughout several evolutionary phases (see, e.g., Pringle \\& Rees 1972, Illarionov \\& Sunyaev 1975): 1) the initial obscured radio pulsar phase in which the pressure of the pulsar radiation is sufficient to keep plasma of the companion stellar wind away from the neutron star magnetosphere; the radio emission of the pulsar is partly (or completely) absorbed by the plasma cloud surrounding the binary system thus the pulsar is practically unobservable; the pulsars spins down due to the magnetodipole radiation; 2) the propeller phase in which the radiation pressure reduced due to the field decay and spin down cannot prevent the infalling matter from an interaction with the magnetosphere but the spin rotation is sufficiently rapid yet to eject the wind matter due to the centrifugal force and to transfer the angular momentum from the neutron star to the wind matter; 3) the wind accretion phase in which the matter from the wind falls on to the surface of the neutron star; nuclear burning of the accreted material heats the neutron star and causes the accretion-driven field decay; in the course of this phase the neutron star can slightly spin up since the accreted matter carries a certain amount of the angular momentum; 4) the enhanced accretion phase which starts when the secondary star leaves the main-sequence and fills its Roche-lobe; a heavy mass transfer on to the neutron star heats the interior to a very high temperature $\\sim 10^{8}-10^{8.5}$K (Fujimoto et al. 1984, Miralda-Escude et al. 1990) and accelerates the accretion-driven field decay; a steady Keplerian disc is formed by the accretion flow outside the neutron star magnetosphere and the accretion torque spins up the neutron star to a very short period. The neutron star processed in the above evolutionary picture can probably prolong its life as a radiopulsar with relatively low magnetic field and short spin period after accretion due to Roche-lobe overflow is exhausted. The properties of pulsars produced due to evolution in low-mass binaries have not been analysed in detail yet and we address this problem in the present paper. Note that a similar problem has been considered by Jahan Miri \\& Bhattacharia (1994) who assumed that the magnetic evolution is determined by the expulsion of the field lines from the neutron star core due to interaction between the fluxoids in the proton superconductor and the vortices in the neutron superfluid. The authors, however, restricted themselves to the case of wide low-mass binaries and considered the evolution only during the main-sequence life time. The paper is organized as follows. We describe in detail the adopted model in Section 2. The results of calculations are presented in Section 3. In Section 4, we discuss our results and compare them with the available observational data on the evolution of neutron stars in low-mass binaries. ", + "conclusions": "We examined the magnetic and spin evolution of neutron stars in close binary systems with a low-mass companion. Due to a long lifetime of the companion, evolutionary transformations of the neutron star may last as long as $10^{9}-10^{10}$ yr and may be very complex. Likely, most of neutron stars in such binary systems experiences a heavy mass transfer ensued due to Roche-lobe overflow. This mass transfer can be extremely extended and the total amount of the mass accreted by the neutron star can exceed $0.1-0.5 M_{\\odot}$ (see, e.g., van den Heuvel \\& Bitzaraki 1995). Besides, the evolution of neutron stars in low-mass binaries has to be influenced by the stellar wind during the main-sequence life of the companion. Due to this, the neutron star in the course of its evolution passes throughout several essentially different evolutionary phases. Calculations presented here are performed in agreement with the standard scenario of evolution of the neutron star in a close binary system (Pringle \\& Rees 1972, Illarionov \\& Sunyaev 1975). According to this scenario, the neutron star can be processed in four main evolutionary phases. Initially, when the star rotates rapidly and the field is strong, the pressure of the pulsar radiation does not allow the wind matter to penetrate into the magnetosphere, and the neutron star evolves like an isolated pulsar. The duration of this stage may vary depending on the initial field strength and the rate of field decay as well as on the intensity of the stellar wind of the companion. The initial phase is obviously shorter for a neutron star with a higher magnetic field and for a stronger wind. Thus, for the initial field as strong as $10^{13}$ G, the phase I may last only $\\sim 0.1-1$ Myr if the neutron star captures gravitationally $\\sim 10^{-12}-10^{-13} M_{\\odot}$/yr from the wind plasma. However, this phase may be extremely long ($\\sim 100-300$ Myr) if the initial field is weak, $B_{0} \\sim 10^{11}$ G or if the rate of gravitational capture is small, $\\sim 10^{-15}-10^{-17} M_{\\odot}$/yr. For the most of considered models, the field is typically reduced by a factor $\\sim 5-10$ during this phase except the cases when its duration is extremely long. When the energy loss due to magnetodipole radiation slows down rotation to such extent that the wind matter can interact with the magnetosphere, the neutron star begins to work as a propeller ejecting the wind plasma. Obviously, the propeller action is more efficient for a stronger magnetic field. The \"propeller\" phase is usually more extended than the phase I. During the \"propeller\" phase, the magnetic evolution of the neutron star in a binary does not differ from that of an isolated star. Since the star is relatively cool to the beginning of the phase II, the conductivity of the crustal matter is determined by impurity scattering and may be high for a low impurity content. Due to this, typically, the field decays less during a more extended \"propeller\" phase than during a shorter \"isolated pulsar\" stage. The spin evolution proceeds, however, more rapidly because the rate of angular momentum loss is larger for the propeller effect than for the magnetodipole braking. During the \"propeller\" phase the neutron star can spin down to a very long period $\\sim 10^{2}-10^{4}$ s. Note that this period may be shorter for stars with low magnetic fields and for a stronger stellar wind. In the further evolution, the spin period of the neutron star can reach the critical value given by equation (11). When rotation becomes so slow, the centrifugal force is unable to eject the wind plasma penetrating into the magnetosphere, and the matter falls on to the neutron star surface. Nuclear burning of the accreted material heats the neutron star, decreases the crustal conductivity and accelerates the field decay. Note that an efficiency of the accretion-driven mechanism of the field decay is strongly sensitive to the accretion rate. The effect of accretion is relatively small for a low accretion rate, $\\dot{M} < 10^{-13} M_{\\odot}$/yr, but it may be of a great importance for higher accretion rates. During the wind accretion, the neutron star slides down the corresponding spin-up line and the rate of this sliding is determined by the rate of the field decay. Depending on the duration of this phase and the rate of accretion, the field can be reduced by a few orders of magnitude. Since the accreted wind matter carries a some amount of the angular momentum, the neutron star can also spin up slightly during this phase but, for the most of considered cases, the spin period is still longer 1 s to the end of the main-sequence evolution of a companion. Usually, the wind accretion phase is the most extended among all evolutionary phases. Note, however, that the star can miss this stage in some cases (low rate of mass loss by a secondary star, rapid decay of the magnetic field), thus the evolution goes directly from the \"propeller\" phase to an enhanced accretion. In low-mass binaries, the neutron star experiences the most dramatic changes during a heavy mass transfer which starts when the companion ends the main-sequence evolution and fills its Roche-lobe. Probably, a mass transfer in such systems can last as long as $10^{7}-10^{8}$ yr. Nuclear burning of the accreted material can heat the neutron star interior to the temperature $T \\sim (2-3) \\times 10^{8}$ K and substantially reduce the crustal conductivity. Due to this, the accretion-driven mechanism of the field decay is especially efficient during the phase IV. Typically, the field may be reduced about 20-30 times if accretion lasts $10^{8}$ yr, but the decrease may be larger for a longer accretion. Depending on the initial magnetic configuration and the duration of the main-sequence evolution, the surface magnetic field can fall down to the value $\\sim 10^{8}- 10^{9}$ G to the end of a mass transfer. Since the matter accreted from the disc carries a large angular momentum, the spin evolves very rapidly. The neutron star reach a spin-up line corresponding to the enhanced accretion on a short time scale, $\\sim 0.01-1$ Myr. In its further evolution, the star moves along the spin-up line while the magnetic field enables to maintain a balance in spin up and the rate of the field decay. During this phase, the accretion torque spins up the neutron star to a very short period, $\\sim 1-100$ ms. Our calculations show that the neutron star with the crustal magnetic configuration can maintain a relatively strong field during the whole evolution in low-mass binaries. Even if the main-sequence life of a companion lasts as long as $10^{10}$ yr and, after that, the neutron star experiences a strong accretion with $\\dot{M} \\sim 10^{-9} M_{\\odot}$/yr during $10^{8}$ yr, the magnetic field can be sufficiently strong to the end of the accretion phase. When accretion is exhausted, such a neutron star works as a radio pulsar with a weak magnetic field, $\\sim 10^{8}-10^{10}$ G, and a short period, $P \\sim 1-100$ ms. These parameters are close to those of millisecond pulsars and, perhaps, these objects are formed in accordance with the considered scenario. Note that if the main-sequence evolution is shorter and a mass transfer is not so extended, the neutron star processed in the above evolutionary transformations should have a stronger magnetic field and a longer period to the end of accretion. This gives a natural explanation of the origin of those radio pulsars in binary systems which lie between millisecond pulsars and the main pulsar population in the B-P plane." + }, + "9802/astro-ph9802278_arXiv.txt": { + "abstract": "Halo white dwarfs can provide important information about the properties and evolution of the galactic halo. In this paper we compute, assuming a standard IMF and updated models of white dwarf cooling, the expected luminosity function, both in luminosity and in visual magnitude, for different star formation rates. We show that a deep enough survey (limiting magnitude $\\gapprox 20$) could provide important information about the halo age and the duration of the formation stage. We also show that the number of white dwarfs produced using the recently proposed biased IMFs cannot represent a large fraction of the halo dark matter if they are constrained by the presently observed luminosity function. Furthermore, we show that a robust determination of the bright portion of the luminosity function can provide strong constraints on the allowable IMF shapes. ", + "introduction": "One way to understand the structure and evolution of the Galaxy is to study the properties of one of its fossil stars: white dwarfs. Their luminosity function has been extensively studied since it provides important information about the properties of the Galaxy and new deep surveys will open the possibility to observe the white dwarf population beyond the cutoff reported by Liebert, Dahn \\& Monet (1988) and Oswalt et al. (1996) as well as to discriminate, on the basis of their kinematical properties, those that belong to the halo. If the halo was formed sometime before the disk as a burst of short duration (Eggen et al. 1962), it would be possible to obtain information about the time elapsed between the formation of both structures (Mochkovitch et al. 1990). If mergers of protogalactic fragments have played an important role in the formation of the galactic halo (Searle \\& Zinn 1978), their signature should be apparent in the white dwarf luminosity function. The observed properties of halo white dwarfs are very scarce. These properties can be summarized as follows: a) Liebert, Dahn \\& Monet (1989) provided a very preliminary luminosity function using six white dwarfs, which were identified as halo members because of their high tangential velocities. b) Flynn, Gould \\& Bahcall (1996) have found that the number of stellar objects in the Hubble Deep Field (HDF) with $V-I>1.8$ is smaller than 3, while M\\'endez et al. (1996) have identified 6 objects with $0 0.3$ at the 90\\% confidence level and $\\Omega_m \\sim 1$ still viable. Blanchard, Bartlett, \\& Sadat (1998) find almost identical results ($\\Omega_m \\sim 0.74$, with $0.3 < \\Omega_m < 1.2$ at the 95\\% confidence level) from these data. Finally, our value of $\\Omega_m$ is inconsistent with the values found by Bahcall, Fan, \\& Cen (1997) ($\\Omega_m = 0.3 \\pm 0.1$), Fan, Bahcall, \\& Cen (1997) ($\\Omega_m \\approx 0.3 \\pm 0.1$), and Bahcall \\& Fan (1998) ($\\Omega_m = 0.2^{+0.3}_{-0.1}$). Our value of $n$ is consistent with the values found by Henry \\& Arnaud (1991) ($n = -1.7^{+0.65}_{-0.35}$) and Henry \\etal (1992) ($n = -2.10^{+0.27}_{-0.15}$), where these authors set $\\Omega_m = 1$. Our value of $n$ is also consistent with the value found by Eke \\etal (1998) ($n = -1.69^{+0.12}_{-0.07}$), where these authors included $\\Omega_m$ as a free parameter. Our value of $n$ is somewhat consistent with the value that Bahcall, Fan, \\& Cen (1997), Fan, Bahcall, \\& Cen (1997), and Bahcall \\& Fan (1998) adopted ($n = -1.4$). Taken as an ensemble, these results are perhaps discouraging in that they span the entire range of acceptable solutions: $0.2 \\la \\Omega_m \\la 1$. This suggests that as yet unknown systematic effects may be plaguing some, if not all, of these results. We briefly identify seven areas where systematic effects could enter ours and similar analyses. (1) The first is the Press-Schechter mass function itself; however, numerical simulations (e.g., Eke \\etal 1996; Bryan \\& Norman 1997; Borgani \\etal 1998) consistently show that the Press-Schechter mass function is an adequate approximation. (2) The spherical collapse model of cluster formation (equations (\\ref{del}) and (\\ref{del2})) may be inadequate. For example, numerical simulations by Governato \\etal (1998) suggest that in equation (\\ref{del2}), the expression $1.69(1+z)$ may really be as low as $\\sim 1.6(1+z)^{0.9}$. This suggests that use of the spherical collapse model may lead to underestimated values of $\\Omega_m$; however, this is only a $\\la 10$\\% effect. (3) Technically, equation (\\ref{virial}) only holds when $\\Omega_m = 1$. Recently, Voit \\& Donahue (1998) derived a virial theorem that holds for all values of $\\Omega_m$, and that allows for the fact that clusters grow gradually; their mass-temperature ($M$-$T$) relation reduces to equation (\\ref{virial}) when $\\Omega_m = 1$. We find that $M$-$T$ relations with functional forms that are similar to that of the $M$-$T$ relation of Voit \\& Donahue (1998) reduce our fitted value of $\\Omega_m$ by $\\la 10$\\%; however, further investigation and use of this $M$-$T$ relation is clearly needed. (4) Also on the subject of the $M$-$T$ relation, care must be taken when fitting to X-ray cluster temperature catalogs: cooling flows lower the measured temperature of most X-ray clusters, which should systematically affect values of $\\Omega_m$ that are determined in this way. (5) Based upon the cooling flow corrected X-ray cluster temperature catalogs of Markevitch (1998) and Allen \\& Fabian (1998), Reichart, Castander, \\& Nichol (1998) determined an empirical $L$-$T$ relation between measured luminosities and cooling flow corrected temperatures that holds for $z \\la 0.5$ and for luminosities that are typical of X-ray cluster catalogs (see \\S3.1); however, more cooling flow corrected X-ray cluster temperature measurements are needed to determine what, if any, exceptions exist to this $L$-$T$ relation, and to extend it to higher redshifts. (6) The art of determining an X-ray cluster catalog's selection function is a constantly improving science; modern selection functions are determined via extensive numerical simulations. An alternative explanation to our high-$\\Omega_m$ result is that the EMSS, for whatever reasons, missed many high-redshift, high-luminosity X-ray clusters beyond what is accounted for by their selection function (see \\S3.3). However, given that the EMSS detected many high-redshift, low-luminosity X-ray clusters, this seems to be an unlikely scenario. (7) Finally, our cosmological model may be inadequate. We did not investigate the effects of a cosmological constant in this paper; however, many authors have demonstrated that the inclusion of a cosmological constant has little effect upon the determined value of $\\Omega_m$ (see \\S1). Also, the effects of quintessent and other exotic cosmologies have not yet been investigated in this context. Some of these potential sources of systematic error can be safeguarded against. For example, the spherical collapse model, the $M$-$T$ relation, and the $L$-$T$ relation all have proportionality factors that are potential sources of systematic error. However, as we have shown in \\S2.1, all of these factors, as well as the parameter $\\sigma_8$, group together, giving us our parameter $c$. Since we fit for $c$, these factors cannot bias our result. However, inadequate functional forms for these relations, as well as for the other functions listed above, can bias ours and others' results. In addition to further theoretical and numerical development of this formalism, only the continued construction of X-ray cluster catalogs will act to further resolve these issues. Fortunately, the number of X-ray cluster luminosity catalogs is growing rapidly. One such catalog is the Southern Serendipitous High-Redshift Archival {\\it ROSAT} Catalog (Southern SHARC) (Collins \\etal 1997; Burke \\etal 1997). The redshift and luminosity ranges of the Southern SHARC are $z < 0.7$ and $L_1 < 3 \\times 10^{44}$ erg s$^{-1}$ (0.5 - 2.0 keV). Although the Southern SHARC does not span the luminosity range of the EMSS, it will provide a good consistency check of our EMSS results. Our analysis of this catalog is underway. Two similar X-ray cluster catalogs that can serve a similar purpose are the {\\it ROSAT} Deep Cluster Survey (RDCS) (Rosati \\etal 1998) and the Wide Angle {\\it ROSAT} Pointed Survey (WARPS) (Jones \\etal 1998). The RDCS spans the redshift and luminosity ranges $z < 0.8$ and $L_1 < 3 \\times 10^{44}$ erg s$^{-1}$ (0.5 - 2.0 keV). The WARPS spans the redshift and luminosity ranges $z < 0.7$ and $L_1 < 2 \\times 10^{44}$ erg s$^{-1}$ (0.5 - 2.0 keV). An analysis of the RDCS is also underway (Borgani \\etal 1998). The 160 deg$^2$ survey (Vikhlinin \\etal 1998b) and the Bright SHARC (Romer \\etal 1998), a high-luminosity extension of the Southern SHARC that is currently under construction, span redshift and luminosity ranges that rival those of the EMSS. Consequently, these catalogs will provide strong, independent checks of the EMSS results. Finally, local ($z_{eff} \\sim 0.1$) X-ray cluster catalogs, such as the {\\it ROSAT} BCS, are of great importance. Although these samples do not have the redshift leverage to constrain cosmological parameters, their large sizes make them excellent samples to better constrain the parameters $n$ and $c$. Samples like the EMSS, the 160 deg$^2$ survey, and the Bright SHARC do not have sufficient luminosity leverage to strongly constrain these parameters, which leads to weaker constraints upon the cosmological parameters. However, a simultaneous analysis of a local X-ray cluster catalog $-$ as opposed to a local X-ray cluster luminosity function as we have used in this paper $-$ and any of these high-redshift, high-luminosity X-ray cluster catalogs could lead to significantly improved constraints upon all of these parameters." + }, + "9802/astro-ph9802223_arXiv.txt": { + "abstract": "s{ Supersymmetric models allow for stable non-topological solitons, Q-balls, which can be produced in the early Universe and contribute to dark matter. Experimental signature of electrically neutral Q-balls is, in fact, the same as is expected for superheavy magnetic monopoles catalyzing baryon decay. Here we use the upper limits on monopole flux obtained with the deep underwater Cherenkov array {\\sf Gyrlyanda} which operated in the Baikal lake in 1984-90 with 267 days of live time to obtain the limit on Q-ball flux. The last has been found to be equal to 3.9 $\\times$ 10$^{-16}$ cm$^{-2}$ sr$^{-1}$ s$^{-1}$ (90$\\%$ CL). This result is discussed and compared with other restrictions. } \\pagebreak ", + "introduction": "Supersymmetric models allow for a new class of objects, which were named Q-balls~\\cite{Coleman85}. These stable non-topological solitons can be produced in the early Universe and contribute to dark matter~\\cite{Coleman85,Kus1,Kus2}. The possible experimental signatures of Q-balls were considered recently by Kusenko, Kuzmin, Shaposhnikov and Tinyakov~\\cite{Kuz}. The interactions of Q-balls with matter were shown in their work to differ essentially on whether they are electrically neutral (SENS) or charged (SECS). Q-balls of SENS type absorb the nuclei with a cross-section $$ \\sigma \\sim 10^{-33}Q^{1/2}(1TeV/m)^{2} cm^{2},\\eqno (1) $$ \\noindent were {\\it m} is assumed to be in the range of 0.1 $\\div$ 100 TeV and {\\it Q} is a soliton charge (baryon number) which must be not less than $\\simeq$ 10$^{15}$({\\it m}/1TeV)$^{4}$ and may be much greater~\\cite{Kus2}. The released energy ($\\sim$ 1 GeV per nucleon) is emitted in pions which (together with their decay products) may become the sources of the Cherenkov radiation in a transparent media. The Coulomb barrier prevents the absorbtion of the incoming nuclei by Q-ball of SECS type, which dissipates its energy in collisions with the matter atoms. In spite of enormous energy released by SECS passing, e.g., the water ($\\sim$ 100 GeV/cm) it does not result in the Cherenkov light. So, SENS passing the water media look very much like the monopoles catalyzing baryon decay~\\cite{ruba} for which the strong experimental flux limits were set with the Baikal deep underwater Cherenkov array {\\sf Gyrlyanda}~\\cite{girl}. In this short note we give the upper limit on Q-ball (SENS) flux which was recalculated from monopole flux limits obtained with {\\sf Gyrlyanda} in 1984-90. Though no experimental restrictions on Q-ball flux have been published by other group so far we compare the {\\sf Gyrlyanda} results with those that can be set using the monopole flux limits. ", + "conclusions": "The upper limit on Q-ball flux of SENS type has been set by revising the old monopole limits obtained by Baikal deep underwater Cherenkov array {\\sf Gyrlyanda}: $F = 3.9 \\times 10^{-16} cm^{-2} sr^{-1} s^{-1}$ (90$\\%$ CL). It is valid for nuclei absorbtion cross-section $\\sigma > 1.9 \\times 10^{-22} cm^{2}$. It is still above the limit obtained for SENS by Baksan telescope but can be improved with the deep underwater neutrino telescope {\\sf NT-200} which is currently under construction in the Baikal Lake. The main advantage of underwater arrays operating in the open water volume and searching for the objects like magnetic monopoles and Q-balls which are expected to generate the intensive light flux passing through a water media is that the effective area is determined mainly by light flux intensity and water optical characteristics (in contrast to underground detectors whose effective area is limited by their geometrical sizes) and may be as large as 10$^{3}$--10$^{5}$ m$^{2}$ even for considerable small arrays." + }, + "9802/astro-ph9802015_arXiv.txt": { + "abstract": "Observations of the Local Arm have been carried out in the near-infrared with the 1.5 m ``Car\\-los S\\'an\\-chez'' telescope in Tenerife. A model of the disc with adjustable parameters fitted to reproduce the DIRBE-COBE survey, has been subtracted from the observational data in order to obtain a clean map of the Local Arm, uncontaminated by other components of the disc. The arm is more than 70 pc over the plane and is wider than 200 pc. At a latitude of about $80^\\circ$ the deviation of the arm from the galactic plane is so large that was only partially observed with our observational window of $\\triangle b = \\pm 6^\\circ$. The elevation over the plane vanishes and the width decreases as the Arm comes closer to the Sun. ", + "introduction": "Because of its proximity, considerable attention has been paid to the spiral feature known as the Local Arm. Optical and 21 cm observations have long been available. Buss et al. (1994) in the ultraviolet, Fatoohi et al. (1996) in gamma-rays, Oliver et al. (1996) in CO, have, more recently, among others, reported observations at other wavelengths. Dynamic properties have been analyzed by Palous (1987), Comeron \\& Torra (1991) and others. In the near infrared, recent surveys (e.g. Odenwald \\& Schwartz 1993; Freudenreich et al. 1994; Garz\\'on et al. 1993) have provided significant information, but the application of these to the study of the Local Arm has been insufficient. Ortiz \\& Lepine (1993) themselves found their description of the Local Arm to be unsatisfactory. More infrared observations are necessary and the main purpose of this work is to compensate for this shortcoming. We report our own observations with the 1.5 ``Carlos S\\'anchez'' telescope in Tenerife, used to search the infrared morphology of the closest spiral feature; this was undertaken bearing in mind the following goals: 1.- To observe the first quadrant region of the arm, which in practice limits attention to the longitude range $70^\\circ-90^\\circ$. Observations extend to slightly higher longitudes, but here interpretation becomes difficult because the arm is too close. 2.- To use high spatial resolution techniques in order to observe individual stars. We determine stellar counts instead of total fluxes. We can in this way perform detailed analysis at various magnitude ranges. The immediate advantage is that this procedure constitutes a natural filter: very bright magnitude stars must, necessarily, be nearby. The importance of the contribution of other components is thus minimized. This probably compensates for the considerable observation time, which in practice means a reduced number of scans. COBE data, for instance, are poorer resolution flux data, and are therefore less suitable to study a close feature like the Local Arm, as they are contaminated by the contribution of larger scale features behind it. 3.- To discriminate between the different contributions to the star counts from other galactic components, mainly the exponential component of the disc. It is to be emphasized that the northern warp of our galaxy is just behind the Local Arm in its first quadrant portion. ", + "conclusions": "In order to better interpret Fig. \\ref{brazo}, we calculated the curve [mean b,l] to quantify the elevation of the Local Arm over the mean plane. This is depicted in Fig.~\\ref{medio}. \\begin{figure} \\resizebox{8.8cm}{!}{\\includegraphics{figu8.ps}} \\caption[]{Mean latitude for the profile of star counts after subtracting the exponential component of the disc. The Local Arm begins at l=$70^\\circ$. \\label{medio}} \\end{figure} As the projected configuration of the arm is reasonably well known from Becker \\& Fenkart (1970) we have also plotted [mean z,l] in Fig. \\ref{zlongitud}. The projected configuration of the arm is a logarithmic spiral with parameters taken from Wainscoat et al. (1992). Furthermore, longitude can be translated into real distance $S$, which is exactly the length along the arc, taking as the origin a point in the arm where $l=70^\\circ$, about 3000 pc away. Figure \\ref{zporbrazo} shows the curve [mean z,$S$]. The arm is elevated about 70 pc and this elevation decreases when the arm comes closer to the Sun, where the elevation practically vanishes. The other part of the arm lies in the direction around $l=270^\\circ$ and cannot be observed in the northern hemisphere. \\begin{figure} \\resizebox{8.8cm}{!}{\\includegraphics{figu9.ps}} \\caption[]{Mean height of the Local Arm in galactic longitude. \\label{zlongitud}} \\end{figure} \\begin{figure} \\resizebox{8.8cm}{!}{\\includegraphics{figu10.ps}} \\caption[]{Mean height of the Local Arm related to the real distance: the length along the arm. \\label{zporbrazo}} \\end{figure} We have also calculated the width of the arm, and Fig. \\ref{ancho} shows the [width,$S$] curve. It must be interpreted with caution: as scans are limited in latitude ($|b| \\leq 6^\\circ$) the arm width and the arm mean position are underestimated. We have seen that the angular width of the arm becomes very large close to the Sun. But when angles are translated to real widths, we see that the arm is thinner as the Sun is approached. At $S$=0 the width is larger than 200 pc. \\begin{figure} \\resizebox{8.8cm}{!}{\\includegraphics{figu11.ps}} \\caption{Mean width of the Local Arm in relation to the distance along the arm.\\label{ancho}} \\end{figure} In Fig. \\ref{brazo} a sharp peak is observed around $l=60^\\circ$. The identification of this feature is not easy but it is clear that it corresponds to a real peak in the number of stars found in that direction, and not to an error in the reduction processes. Such a peak is not observed in the DIRBE data. Possible explanations should take into account the patchy nature of the absorbing clouds, and the different resolution of DIRBE and our data, but we have not found a relation with any other known feature of the galaxy, nor with detectable especial features in CO or 21 cm maps. At a galactic latitude around $80^\\circ$, a large and unexpected deviation of the arm with respect to the mean plane prevents us to obtain reasonable values of the mean position of the Local Arm, so that the values in our Fig. 8 are highly underestimated. The Arm is to a great extent outside the small observation window of $\\triangle b = \\pm 6^\\circ$. Clearly, future observations should be carried out for $|b| \\leq 12 ^\\circ$ in order to assure better observing conditions and better quantitative results. Nevertheless, even if incomplete, our map in figure 8 constitutes a first observation at this wavelength. This fact together with the substraction of the foreground sources which has been carried out, makes our observations an important source of information to study the closest known spiral arm. There is a very noticeable agreement between our [mean z,l] curve and the similar one obtained by Kolesnick \\& Vedenicheva (1978). This agreement is mainly due to the fact that we are actually observing young stars. Near infrared surveys are usually used to trace the distribution of old stars, but due to the proximity of the Local Arm, what we are mainly observing are OB stars with $M_K \\approx -4.7$. Nevertheless observations at different wavelengths must be obtained even if they confirm the optical description. Our map, on the other hand, is uncontaminated by other sources, i.e. by the exponential component of the disc, warp included. This is important, as the positions of the arm and the north warp maximum are found in the same direction. Figure \\ref{forma} represents the geometrical configuration of the Local Arm, when viewed from approximately the galactic centre. \\begin{figure} \\resizebox{8.8cm}{!}{\\includegraphics{figu12.ps}} \\caption{Geometrical configuration of the Local Arm viewed from the galactic centre. \\label{forma}} \\end{figure}" + }, + "9802/astro-ph9802001_arXiv.txt": { + "abstract": "We consider, herein, a model for $\\gamma$-ray production in blazars in which a relativistic, highly-collimated electron-proton beam interacts with a dense, compact cloud as the jet propagates through the broad and perhaps narrow line regions (BLR and NLR) of active galactic nuclei (AGN). During the propagation of the beam through the cloud, the process of excitation of plasma waves becomes an important energy loss mechanism, especially for mildly relativistic proton beams. We compute the expected spectra of $\\gamma$-rays from the decay of neutral pions produced in hadronic collisions of the beam with the cloud, taking into account collisionless losses of the electron-proton beam. This model may explain the X-ray and TeV $\\gamma$-ray (both low and high emission states) of Mrk 421 as a result of synchrotron emission of secondary pairs from the decay of charged pions and $\\gamma$-ray emission from the decay of neutral pions for the plausible cloud parameters. However clouds can not be too hot and too dense. Otherwise the TeV $\\gamma$-rays can be attenuated by the bremsstrahlung radiation in the cloud and the secondary pairs are not able to efficiently produce synchrotron flares because of the dominant role of inverse Compton scattering. The non-variable $\\gamma$-ray emission observed from Mrk 421 in the EGRET energy range cannot be described by the $\\gamma$-rays from decay of neutral pions provided that the spectrum of protons in the beam is well described by a simple power law. These $\\gamma$-rays might only be produced by secondary pairs scattering the soft non-variable X-rays which might originate in the inner part of the accretion disk. ", + "introduction": "Many blazar-type active galaxies have recently been detected in MeV - GeV $\\gamma$-rays by detectors on the board of Compton GRO (von Montigny et al.~1995). Mrk 421, Mrk 501, and 1E2344 are observed in the TeV $\\gamma$-rays (Punch et al.~1992, Quinn et al.~1996, Catanese et al.~1997a). The $\\gamma$-ray emission of blazars is highly variable on different time scales (from months and weeks up to a fraction of an hour, e.g. Gaidos et al. 1996, Mattox et al. 1997). The spectra of FSRQ blazars are well described by a power-law with a spectral break from the lower energy part observed in a few sources (McNaron-Brawn et al. 1995, Sch\\\"onfelder et al. 1996), and possibly a cut-off at higher energies (Pohl et al. 1997). The spectra of two BL Lacs, Mrk 421 and Mrk 501, extend up to at least $\\sim 10$ TeV (Krennrich et al. 1997, Aharonian et al. 1997). There is also weak evidence of emission of $\\sim 50$ TeV photons from some nearby BL Lacs (Mayer \\& Westerhoff 1996). These observations are usually interpreted in terms of the inverse Compton scattering model. This model has been discussed previously in many different geometrical scenarios (see for example reviews by Dermer \\& Schlickeiser 1992, Sikora 1994, Dermer \\& Gehrels 1995, Schlickeiser 1996, Bednarek 1998). The $\\gamma$-rays can be also produced in collisions of very high-energy hadrons with soft radiation (e.g. Sikora et al. 1987). This mechanism is discussed in the context of $\\gamma$-ray production in blazars by e.g. Mannheim \\& Biermann (1992), Coppi, Kartje, K\\\"onigl~(1993), Protheroe (1997), Bednarek \\& Protheroe (1997). The interaction of hadronic jets with background matter in AGNs as a possible source of $\\gamma$-rays has been problematic because of the difficulty of finding a dense enough target. In principle, such a target might be created either by the matter of a thin accretion disk (Nellen et al. 1993) or by the matter in a corona of a thick accretion disk (Bednarek 1993). The $\\gamma$-ray emission observed from blazars might originate in hadronic collisions of a highly collimated proton beam with clouds entering the jet. This model, recently proposed by Dar \\& Laor (1997), is similar to the scenario developed by Rose et al.~(1984, 1987). In the broad line region (BLR), the typical dimensions of clouds are of order $\\sim 10^{12-13}$ cm with cloud densities $\\sim 10^{10-12}$ cm$^{-3}$. For these parameters, the hadronic collisions occur frequently enough to produce the observed fluxes of $\\gamma$-rays. However, collisionless processes principally driven by the two-stream instability (Scott, et al.~1980, Rose et al.~1984, 1987) can be the dominant energy loss mechanism for such jets in parameter ranges associated with the BLR and NLR (Beall 1990). The computations by Dar and Laor do not take into account these collisionless energy losses as the jet propagates through the dense, hot medium of the cloud. Where collisionless energy losses are significant, propagation lengths are markedly shortened. This is especially true for electron-proton beams with relatively low Lorentz factors (see Rose et al. 1984, and Beall 1990). These effects must be taken into account when calculating the $\\gamma$-ray emission from the jet. The radiation emitted by a relativistic electron-proton beam interacting with an ambient medium such as a broad line cloud has been previously discussed in general terms by Rose et al.~(1987) and Beall et al. (1987). In this paper the predictions of the model wherein an electron-proton beam interacts with an interstellar cloud are compared with the results of $\\gamma$-ray observations of Mrk 421 in a high and low states. We extend the results of the recent discussion of this model by Dar \\& Laor (1997) to take into account the influence of collisionless losses on the electron-proton jet's $\\gamma$-rays spectra. ", + "conclusions": "If the relativistic electron-proton beam collides with a dense, compact cloud, the energy losses of the jet via the excitation of plasma waves (Rose et al 1984) becomes important for beams with energies below $\\sim 300$ GeV (Fig.~1). This process causes the break in the spectrum of $\\gamma$-rays from the decay of neutral pions that are produced in inelastic collisions of protons with matter. Such a model can naturally explain the low and high states of TeV $\\gamma$-ray emission in BL Lac type blazars, provided that the proton beam interacts with different column densities of matter. Moreover, because of the collisionless energy losses of the proton beam, the $\\gamma$-ray spectrum predicted in the 100 MeV - 1 GeV energy range does not vary strongly. This features fits nicely to the general behavior of the $\\gamma$-ray emission from Mrk 421 between low and high states. However the $\\gamma$-ray spectra computed in the 100 MeV - 1 GeV energy range for reasonable sets of model parameters are too flat, showing significant deficiency in comparison to the Mrk 421 spectrum observed by EGRET (Fig.~3) provided that the spectrum of relativistic protons in the beam is well described by a single power law. Secondary $e^\\pm$ pairs are also produced in hadronic collisions through decay of charged pions, with the spectra similar to the $\\gamma$-ray spectra but shifted to lower energies. The higher energy, variable part of the spectrum of secondary pairs (with energies above $\\sim 10$ GeV) can be responsible for the simultaneous x-ray flares. Such a picture could work provided that the clouds are not too hot. Otherwise the TeV $\\gamma$-rays are absorbed by the thermal bremsstrahlung photons produced in the cloud (see computations of the $\\gamma$-ray mean free paths on Fig.~3 in Beall et al. 1987). The clouds also cannot be too dense. By reversing Eq.~(9), we see that if the cloud density is \\begin{eqnarray} n_c > 3\\times 10^{25} \\rho_{ph} T_c^{-0.5} \\lambda^{-1}, \\end{eqnarray} \\noindent where $\\lambda = n_c r_c$ is the column density of matter traversed by the proton beam in the cloud, then the secondary electrons are not able to lose energy efficiently by the synchrotron process because of dominant role of ICS losses on thermal bremsstrahlung radiation produced in the cloud. The relation of the break in the spectrum of secondary pairs to the break in the x-ray spectrum observed in Mrk 421 allows a derivation of the strength of the magnetic field in the emission region equal to $\\sim 7.1$ G (Eq.~3). These pairs cannot contribute to the EGRET energy range: by production of $\\gamma$-rays in a bremsstrahlung process; or by inverse-Compton scattering of soft photons of synchrotron origin, or coming from the optically thick, low temperature accretion disk, or bremsstrahlung photons produced in the cloud. Only the scattering of non-variable, soft x-ray photons ($0.1 - 1$ keV) by secondary pairs with energies below $\\sim 10$ GeV could explain the non-variable $\\gamma$-ray emission from Mrk 421 in the EGRET energy range. These photons could be produced in the inner part of the accretion disk." + }, + "9802/astro-ph9802237_arXiv.txt": { + "abstract": "We have examined 8004 plates at the Harvard College Observatory Plate Collection searching for optical transient emission from the gamma-ray burst GRB 970228. This is the first archival search carried out so far for a gamma-ray burst with known transient optical emission. The total exposure time amounts to $\\sim 1.1$ yr. No convincing optical activity was found above $12.5$ mag at the expected position of the GRB 970228 optical counterpart. ", + "introduction": "GRB 970228 was detected as a 5.5-s brief, high energy event on 28 February 1997, by the satellite BeppoSAX (Costa et al. 1997a). A decaying X-ray emission (the X-ray afterglow) was found 8 hr after the onset of the event (Costa et al. 1997b). Optical observations started 15.4 hr after the burst (Pedichini et al. 1997), which allowed the search of variable objects in the intersection between the $\\gamma$-ray error box, the X-ray error box and the Interplanetary Network (IPN) annulus (Hurley et al. 1997). They led to the discovery of the first optical counterpart of a GRB (van Paradijs et al. 1997). The recent measurement of the redshift of the GRB 970508 optical counterpart has established that most GRBs - if not all - lie at cosmological distances (Metzger et al. 1997). While galactic GRB models allow multiple outbursts from each source (Li and Dermer 1992), the cosmological models are usually based on singular, cataclysmic highly energetic events. Thus, the detection of recurrent optical emission for a GRB would impose severe constraints on the cosmological models (Narayan, Paczy\\'nki and Piran 1992, M\\'esz\\'aros and Rees 1993). Archival searches can be crucial in order to clarify this problem. Archival searches in GRB error boxes rely upon the assumption that GRBs do repeat and show optical transient emission. With perhaps the exception of OT 1928 (Schaefer 1981, Hudec et al. 1994), no firm candidates have yet been established (Hudec et al. 1993, and references therein, also Gorosabel and Castro-Tirado 1998). The detection of the two first optical counterparts for GRB 970228 (van Paradijs et al. 1997, Guarnieri et al. 1997) and GRB 970508 (Bond 1997,Castro-Tirado et al. 1998) has been an important breakthrough. However, for some other GRBs no optical emission was detected although very fast and deep follow up observations were carried out (Castro-Tirado et al. 1997, Groot et al. 1998). ", + "conclusions": "The $\\sim$ 8000 plates span $\\sim$1.1 yr. Many star-like spots were found. With the exception of the above-mentioned case, none of them was consistent with the position of~~the~~true GRB 970228 optical counterpart. Therefore we can settle a lower limit of $\\sim$ 1.1 yr for any recurrent optical transient emission activity brighter than 12.5 mag (or 4.2 mag if a 1-s flash is assumed). According to extrapolation of the gamma-ray spectra into longer wavelengths for the strongest bursts observed by BATSE, the magnitude of the optical flash that -eventually -would arise simultaneously to the gamma-ray event, could reach a magnitude of 2.5 in 1-s (for the spectral index $\\alpha$ $\\geq$ --1, Ford and Band, 1996). If this would have been the case for a previous bursting activity at the GRB 970228 location, such an optical flash would have been easily detected in any of the HCO plates." + }, + "9802/astro-ph9802147_arXiv.txt": { + "abstract": "We present the results of nitrogen and oxygen abundance measurements for 185 \\ion{H}{2} regions spanning a range of radius in 13 spiral galaxies. As expected, the nitrogen--to--oxygen ratio increases linearly with the oxygen abundance for high metallicity \\ion{H}{2} regions, indicating that nitrogen is {\\it predominately} a secondary element. However, the nitrogen--to--oxygen ratio plateaus for oxygen abundances less than 1/3 solar (12+log(O/H) $<$ 8.45), as is also seen in low metallicity dwarf galaxies. This result suggests that the observed trend in dwarf galaxies is not due to outflow of enriched material in a shallow gravitational potential. While the effects of infall of pristine material and delayed nitrogen delivery are still unconstrained, nitrogen does appear to have both a primary and secondary component at low metallicities in all types of galaxies. ", + "introduction": "Abundance measurements provide probes of the elemental enrichment process and thus are important tracers of galaxy formation and evolution. Recently, chemodynamical models have attempted to obtain self--consistent star formation histories for galaxies by combining both chemical and dynamical processes (e.g., Samland, Hensler, \\& Theis \\markcite{SHT97}1997). These models require an understanding of the star formation and enrichment processes, including stellar nucleosynthesis, dust depletion, and gas dynamics. While stellar yields for many elements, such as O, Ne, S, and Ar (e.g., Woosley \\& Weaver \\markcite{WW95}1995), are in good agreement with their observed enrichment ratios (e.g., Thuan, Izotov, \\& Lipovetsky \\markcite{TIL95}1995), the effective yields and origin of some elements, such as nitrogen, are still a matter of debate. Understanding the origin of nitrogen is particularly important since it is frequently used to derive the primordial helium abundance via regression analysis of the observed enrichment of helium and nitrogen (e.g., Pagel \\& Kazlauskas \\markcite{PK92}1992). In high metallicity environments, nitrogen is believed to be synthesized during the CNO cycle in intermediate mass stars via a secondary process (Renzini \\& Voli \\markcite{RV81}1981). On the other hand, studies of the N/O ratio in low metallicity \\ion{H}{2} regions suggest that there may also be a primary origin for nitrogen (e.g., Edmunds \\& Pagel \\markcite{EP78}1978; Garnett \\markcite{G90}1990; Thuan \\etal \\markcite{TIL95}1995). In this paper, we examine the relative enrichment of nitrogen and oxygen in spiral galaxy \\ion{H}{2} regions in order to investigate the origin of nitrogen in a wide range of metallicity environments. Previous studies of chemically enriched extragalactic \\ion{H}{2} regions indicate that the effective yield of nitrogen depends on the previous enrichment of C and O, and thus nitrogen is considered a ``secondary'' element (e.g., Torres--Peimbert, Peimbert, \\& Fierro \\markcite{TPPF89}1989; Vila--Costas \\& Edmunds \\markcite{VE93}1993). The signature for secondary nitrogen production is that the N/O ratio increases linearly with O/H. In contrast, studies of the nitrogen enrichment in low mass dwarf galaxies indicate that the N/O ratio is constant for oxygen abundances less than 1/3 solar (e.g., Garnett \\markcite{G90}1990; Thuan \\etal \\markcite{TIL95}1995; van Zee, Haynes, \\& Salzer \\markcite{vHS97b}1997b). A constant N/O ratio may be indicative of primary nitrogen production. Primary nitrogen is defined as nitrogen produced only out of the original hydrogen and helium in a star, either directly or through successive stages of burning. Low metallicity intermediate mass stars which undergo successive dredge--ups of their enriched core are one source of primary nitrogen (Renzini \\& Voli \\markcite{RV81}1981). Recent work on the nucleosynthesis of massive stars suggests that primary N may also be produced in low metallicity massive stars via convective overshoot (Woosley \\& Weaver \\markcite{WW95}1995). However, the signature of primary nitrogen production, a constant N/O ratio at low metallicities, can also be caused by dynamic processes, such as gas infall or outflow. Standard ``closed box'' chemical enrichment models implicitly assume that enriched materials will be retained by the system; this assumption may not be valid in low mass systems where the gravitational potential well is quite shallow. In such objects, the products of nucleosynthesis may be ejected from the system during supernovae explosions. While Edmunds \\markcite{E90}(1990) states that ``simple and pure'' outflows will not affect relative enrichment ratios, such as N/O, gas flows are likely to be neither ``simple'' nor ``pure''. For instance, oxygen is predominately made in high mass stars which undergo more violent deaths, and thus is more likely to be removed from the galaxy. This {\\it differential} outflow results in a decrease in the effective yield for oxygen with a corresponding increase in the N/O ratio. Furthermore, additional complications arise from possible accretion of pristine gas from external reservoirs. This is particularly a concern for dwarf galaxies since many have \\ion{H}{1} envelopes which extend well beyond the optical system (e.g., van Zee, Haynes, \\& Giovanelli \\markcite{vHG95}1995). In dwarf galaxies, the elemental enrichment could be diluted by accretion, resulting in lower effective yields. The effect of gas infall on the N/O ratio is complex and highly model dependent. For primary nitrogen, the N/O ratio will be uneffected by dilution; for secondary nitrogen, the inflow of pristine material can have a wide range of effects on the N/O ratio (Edmunds \\markcite{E90}1990). Most studies on the origin of nitrogen have concentrated on low luminosity dwarf galaxies, where the oxygen abundance is very low. However, these low mass systems are particularly susceptible to both outflow (due to their shallow gravitational potential) and inflow (from extensive neutral gas envelopes). We have taken advantage of the expected low abundance nature of outlying \\ion{H}{2} regions in spiral galaxies (e.g., Zaritsky, Kennicutt, \\& Huchra \\markcite{ZKH94}1994) to investigate whether a large contribution of primary nitrogen is evident in \\ion{H}{2} regions associated with higher mass objects. Previous observations of \\ion{H}{2} regions in spirals (e.g., Vila--Costas \\& Edmunds \\markcite{VE93}1993; Thurston, Edmunds, \\& Henry \\markcite{TEH96}1996) have focussed on the brighter \\ion{H}{2} regions in the inner galaxy, which are correspondingly more enriched. By focussing on {\\it outlying} \\ion{H}{2} regions, we have observed spiral galaxy \\ion{H}{2} regions with metallicities as low as 1/10 solar, comparable to typical dwarf galaxy metallicities. By observing massive spiral galaxies, we expect the complicating effects of gas outflow/infall to be minimized. ", + "conclusions": "" + }, + "9802/astro-ph9802158_arXiv.txt": { + "abstract": "New optical and near-UV HST observations of Geminga are presented. When compared with previous ground-based and HST imaging, the data confirm and better define the presence of an emission feature centered at $\\sim 6,000 \\AA$ and superimposed on the thermal continuum best fitting the extreme-UV/soft X-ray data. This feature may be interpreted in terms of cyclotron emission originated from a mixture of H/He ions in the neutron star's atmosphere. In the case of pure Hydrogen, the feature wavelength would imply a magnetic field of order $3-5~10^{11}~G$, consistent with the value deduced from the dynamical parameters of the pulsar. If due to a cyclotron emission, the observation of this feature would represent the first case of an {\\it in situ} measurement of the surface magnetic field of an isolated neutron star. ", + "introduction": "Our understanding of the optical properties of Isolated Neutron Stars (INSs) is limited by the faintness of the vast majority of them. Indeed, only for the $m_{V} \\sim 16.6$ Crab pulsar, acceptable, medium-resolution optical spectroscopy is available (Nasuti et al 1996). For few more cases (PSR0540-69, PSR0833-45, PSR1509-58, PSR0656+14 and Geminga), the spectral information just relies on multicolour photometry (Nasuti et al. 1997; Mignani et al. 1998a; Pavlov et al. 1997; Bignami et al. 1996). For the rest of the optical database (PSR0950+08, PSR1929+10 and PSR1055-52), only one-or two-band detections are available (Pavlov et al. 1996; Mignani et al. 1997). \\\\ Very young objects, say up to $\\tau \\sim 10^{4}$ yrs (Crab, PSR0540-69, PSR0833-45), are characterized by flat, synchrotron-like spectra arising from energetic electron interactions in their magnetosphere (Caraveo, 1998). For the Middle-Aged ($\\tau \\sim 10^{5}~ yrs$) INSs (PSR0656+14, Geminga and PSR1055-52), also referred to as \"the three musketeers\" for their overall similarities (Becker \\& Tr\\\"umper 1997), the situation is more complex and different emission processes may become relevant. For example, the non-thermal magnetospheric emission could have faded enough, at least in the optical waveband, to render visible the thermal emission from the hot neutron star surface (Mignani et al. 1998b). Its temperature, following standard cooling calculations (see e.g. Nomoto \\& Tsuruta 1987), could be in the $10^{5}- 10^{6} ~ K$ range, in excellent agreement with the recent X-ray findings (e.g. Becker \\& Tr\\\"umper 1997). It is easy, then, to predict the IR-optical-UV fluxes expected along the $\\sim E^{2}$ Rayleigh-Jeans slope of the Planck curve best fitting the X-rays, and to compare predictions with observations, where available. ", + "conclusions": "The new data presented here, improving as they do the data given in Bignami et al. (1996), leave no room for doubt that a wide emission feature exists around $6,000 \\AA$, superimposed on a possibly thermal continuum. The feature falls in the wavelength region expected for an atmospheric proton-cyclotron emission produced close to the surface of a magnetic neutron star. Geminga is a magnetic neutron star, to wit its periodic $\\gamma$-ray emission (Mattox et al, 1998), and most probably has an atmosphere, to wit the spectral shape of its soft X-ray emission (Meyer et al 1994, Halpern \\& Rudermann 1993). \\\\ Not only the optical/UV data may represent, historically, the first time we actually see the surface (atmosphere) of an INS: our interpretation of the feature reported here would also imply the first direct measure of the magnetic field of an INS.\\\\ The question arises wheter to expect similar features for other INSs, as the X-ray/UV/optical data improve. If the physical assumptions mentioned above are correct, then, by and large, the answer is positive. We do expect to observe similar features on other INSs of comparable age, first of all on the other two \"musketeers\". However, for PSR0656+14 the photometry information available so far (Pavlov et al. 1997) shows a rather different spectral behaviour and the optical flux appears more consistent with a steep power law ($\\alpha \\sim 1.4$), although an additional blackbody component is present. Its optical emission would thus be the combination of both magnetospheric and thermal processes. The recent detection of optical pulsations from PSR0656+14 (Shearer et al. 1997b) seems also to support the presence of a strong magnetospheric component. For the third and oldest musketeer, our recent detection (Mignani et al. 1997) falls close enough to the Rayleigh-Jeans extrapolation of the ROSAT spectrum. \\\\ The difficulties in understanding INS physical properties have been outlined dramatically, even for the three best known cases of X-ray thermal emitters. Of them, Geminga is certainly the one with the most abundant phenomenology at yet, disentangling its thermal and non-thermal emission (i.e. distiguishing surface from magnetosphere), although well under way, will require more refined observations." + }, + "9802/astro-ph9802191_arXiv.txt": { + "abstract": "We present a series of 2-dimensional hydrodynamic simulations of massive disks around protostars. We simulate the same physical problem using both a `Piecewise Parabolic Method' (PPM) code and a `Smoothed Particle Hydrodynamic' (SPH) code, and analyze their differences. The disks studied here range in mass from $0.05 M_*$ to $1.0 M_*$ and in initial minimum Toomre $Q$ value from $1.1$ to $3.0$. We adopt simple power laws for the initial density and temperature in the disk with an isothermal ($\\gamma=1$) equation of state. The disks are locally isothermal. We allow the central star to move freely in response to growing perturbations. The simulations using each code are compared to discover differences due to error in the methods used. For this problem, the strengths of the codes overlap only in a limited fashion, but similarities exist in their predictions, including spiral arm pattern speeds and morphological features. Our results represent limiting cases (i.e. systems evolved isothermally) rather than true physical systems. Disks become active from the inner regions outward. From the earliest times, their evolution is a strongly dynamic process rather than a smooth progression toward eventual nonlinear behavior. Processes that occur in both the extreme inner and outer radial regions affect the growth of instabilities over the entire disk. Effects important for the global morphology of the system can originate at quite small distances from the star. We calculate approximate growth rates for the spiral patterns; the one-armed ($m=1$) spiral arm is not the fastest growing pattern of most disks. Nonetheless, it plays a significant role due to factors which can excite it more quickly than other patterns. A marked change in the character of spiral structure occurs with varying disk mass. Low mass disks form filimentary spiral structures with many arms while high mass disks form grand design spiral structures with few arms. In our SPH simulations, disks with initial minimum $Q=1.5$ or lower break up into proto-binary or proto-planetary clumps. However, these simulations cannot follow the physics important for the flow and must be terminated before the system has completely evolved. At their termination, PPM simulations with similar initial conditions show uneven mass distributions within spiral arms, suggesting that clumping behavior might result if they were carried further. Simulations of tori, for which SPH and PPM are directly comparable, do show clumping in both codes. Concern that the point-like nature of SPH exaggerates clumping, that our representation of the gravitational potential in PPM is too coarse, and that our physics assumptions are too simple, suggest caution in interpretation of the clumping in both the disk and torus simulations. ", + "introduction": "Over the past several years a broad paradigm of star formation has been developed (see Shu, Adams \\& Lizano 1987). First, a cloud of gas and dust collapses and forms a protostar with a surrounding disk. Later the star/disk system ejects matter in outflows as well as continuing to accrete matter from the cloud. Finally, accretion and outflow cease and the star gradually loses its disk and evolves onto the main sequence. While this paradigm provides for a good qualitative picture of the star formation process, many important issues require further work. For example, observations by several groups (Simon \\etal 1995, Ghez \\etal 1993, Leinert \\etal 1993, Reipurth \\& Zinnecker 1993) show that young stars in many different star forming regions are commonly found in binary or higher order multiple systems, with a broad peak in separation distance at around 30 AU. In addition, many of the higher order multiples show hierarchical characteristics: a distant companion orbiting a close binary for example. In what manner are multiple systems such as these formed? A variety of studies have been undertaken to model the processes leading to the observed systems. One class of models begins with the collapse of a cloud of matter. These results (Bate \\etal 1995, Foster \\& Boss 1996, Boss 1995, Burkert \\& Bodenheimer 1993, Bonnell \\& Bastien 1992, Myhill \\& Kaula 1992) show that both single stars and multiple systems can be formed from the collapse and subsequent fragmentation of rotating, spherical or elongated molecular cloud cores. This class of simulations focus on the collapse phase but do not follow in detail the dynamics of disks formed from the material with initially higher angular momentum. In addition, a number of models extended beyond the initial collapse (Bonnell 1994, Pickett \\etal 1996, Woodward \\etal 1994) have shown that post-collapse objects can be driven into fragmentation, or into spiral arm and bar formation prior to the development of a Keplerian disk. Laughlin \\& Bodenheimer (1994) have simulated the evolution of a collapsing cloud in 2D and then followed its late time behavior with a 3D disk simulation. They have found that such a collapse leads to a core plus a long lived, broad torus. The development of $m=1$ and $m=2$ spiral patterns may lead to late time fragmentation of the torus ($m$ is the number of arms in the spiral pattern). As a star-disk or multiple-star-disk system evolves, the dynamics of the disk itself as well as its interaction with the star or binary becomes important in determining the final configuration of the system. Depending on its mass and temperature, a disk may develop spiral density waves and viscous phenomena of varying importance. Each may be capable of processing matter through the disk as well as influencing how the disk eventually decays away as the star evolves onto the main sequence. A variety of mechanisms for production of spiral instabilities in disks around single stars have been suggested. An incomplete list includes the linear perturbation results of Adams, Ruden \\& Shu (1989) (hereafter ARS) who suggest a mechanism (`SLING'-- see Shu \\etal 1990) by which a resonance between the star and a one-armed ($m$=1) spiral mode may become globally unstable. Both perturbation theory (Papaloizou \\& Lin 1989) and numerical calculation (Papaloizou \\& Savonjie 1991, Heemskirk \\etal 1992) have shown another instability mechanism based on the distribution of specific vorticity (termed ``vortensity'') which can influence evolution in disks and tori. It is driven primarily by wave interactions at corotation and can act either to suppress or amplify spiral waves in the disk, depending on the vortensity gradient there. Another family of instabilities is based upon vortensity gradients at the boundaries of the disk or torus. The SWING amplifier (Goldreich \\& Lynden-Bell 1965, Julian \\& Toomre 1966, Goldreich \\& Tremaine 1978) provides an instability channel whereby low amplitude leading spiral arms unwind and are transformed into much larger amplitude trailing waves. A feedback cycle then creates additional leading waves and the instability grows. This paper is a continuation of work by two of us (Adams \\& Benz 1992, hereafter AB92), who began modeling of disks of mass $M_D\\gtrsim 0.5 M_*$ and observed formation of spiral arms and clumps. We present a series of two dimensional numerical simulations of circumstellar disks with masses between $0.05 M_*$ and $1.0 M_*$. We attempt to characterize the growth of instabilities and pay particular attention to the existence and effect of the SLING instability. In section \\ref{codes}, we outline the numerical methods used and discuss the limitations of each code and their effects on our simulations. In section \\ref{phyasmpt}, we outline the initial conditions adopted for the disks studied and in section \\ref{results}, we first describe qualitatively the results of our simulations and then begin a quantitative analysis of the pattern growth, the correspondence between two hydrodynamic codes, and the correspondence between linear analyses and hydrodynamic simulations. In Section \\ref{summary}, we summarize the results and their significance in the evolution of stars and star systems. ", + "conclusions": "\\label{summary} By using two conceptually different hydrodynamic methods (SPH and PPM), we are able to simulate a broader range of problems, but gain a sobering insight into the limitations of these tools. It is striking that PPM indicates violent behavior near the inner boundary (weakly supported by SPH), and that SPH indicates pronounced clumping (weakly supported by PPM). Both methods indicate that instability growth is not a steady progression from low to high amplitude perturbations with a single dominant pattern present throughout. Both methods indicate a marked change in the character of instabilities with disk mass. Low mass disks form many armed filimentary spiral structures while high mass disks form few armed grand design spiral structures. In this study of the evolution of circumstellar accretion disks, we have found simultaneous growth of global spiral instabilities with multiple Fourier components. Growth of each of the components occurs over the course of a few orbit periods of the disk and a single component rarely dominates the evolution of a disk. As expected, the massive disks are found to be the most unstable, due to self-gravitating instabilities within the disk. Accretion of matter onto the star itself can, in warm disks (i.e. those with high \\qmin~values), significantly drain matter from the disk on similar time scales to the self-gravitating instabilities. Short-term variations in the amplitude of a given component, and strong constructive interference behavior between different components, can produce `spikes' in the surface density. These spikes can eventually grow to such amplitude that gravitational collapse occurs resulting in the production of one or more clumps. Pattern growth is stimulated at early times by the rapid growth of instabilities at small radii which eventually engulf the entire disk. Steady spiral arm structures are not generally present. Instead, spiral arms progressively grow, fragment and reform as time progresses. In cases where accretion is rapid, power can be produced in an $m=1$ spiral pattern due to nonaxisymmetric accretion of mass and momentum onto the star. Understanding the dynamics of the inner region is of primary importance for understanding the global morphology of the system. The gross structure of low and high mass disks are markedly different from each other. High mass disks form large, grand design spiral arms with few arms, while low mass disks form predominantly thin, filamentary multi-armed structures. In almost no case is the $m=1$ spiral pattern the fastest growing pattern in the disk. Typically a combination of $m=2-4$ patterns in high mass disks or very high order patterns ($m\\gtrsim 5$) in low mass disks dominate the morphology. The transition between these behaviors comes at approximately \\mrat~$=0.2-0.4$. This transition corresponds to the `maximum solar nebula' mass discussed in STAR, above which $m=1$ modes due to SLING are expected to grow strongly. It is intriguing to speculate that the collapse processes seen here are responsible for the production of brown dwarf-like companions such as that seen by Nakajima \\etal (1995) and/or of planetary companions similar to those recently discovered around several nearby stars (Mayor \\& Queloz 1995, Marcy \\& Butler 1996, Gatewood 1996). However, we must emphasize that clump formation in self-gravitating circumstellar disks depends on the ability of the gas to cool efficiently. Our simulations here use a simple isothermal equation of state which favors clump formation. Additional simulations with realistic cooling functions, including radiative transfer effects, must be done in order to clarify this important issue." + }, + "9802/astro-ph9802228_arXiv.txt": { + "abstract": "Sco X-1 is the brightest extra-solar point source of X-rays, and may serve as a prototype for low mass X-ray binaries as a class. It has been suggested that the UV and optical emission arise as a result of reprocessing of X-rays, and that a likely site for such reprocessing is an accretion disk around the X-ray source. If UV and optical emission are enhanced by reprocessing of X-rays, the X-ray variability may be manifest in UV emission: we test this by using high temporal resolution UV data obtained simultaneously with high temporal resolution X-ray data obtained simultaneously with the GHRS on the Hubble Space Telescope, and with the X-ray Timing Explorer (XTE). We analyze the variability behavior of the UV spectrum and of the X-rays, and we also measure the properties of the emission line profiles as viewed at high resolution (resolving power$\\simeq$25,000) with the echelle gratings. The variability behavior does not provide direct support of the reprocessing scenario, although the correlated variability between UV and X-rays does not conflict with this hypothesis. Furthermore, the emission line profiles do not fit with simple models for disk emission lines. ", + "introduction": "X-ray binaries remain among the most fascinating astronomical systems. Pronounced periodic and aperiodic variability of the X-ray emission, on timescales ranging from milliseconds to decades, distinguishes the X-ray binaries from all other astronomical sources. For example, quasi-periodic oscillations (QPOs) in the X-rays from low mass X-ray binaries (LMXB's) (Lewin {\\it et al} 1988; van der Klis, 1989) may help us understand accretion disks about compact objects near their Eddington luminosities, disk interactions with the magnetic fields of neutron stars, and pulsar evolution (e.g., Lamb 1988). Yet the current X-ray observations lack the ability to measure local physical conditions and dynamics with high-resolution spectrometry. In the optical and UV wavebands we can measure variability of the UV spectra of these sources on short timescales and at spectral resolution adequate to dissect accretion flows according to velocity. Thus, time-resolved UV spectrometry constitutes an important discovery space for these sources. A prime target for such observations is Sco~X-1, the first discovered and brightest extra-solar X-ray source and also the strongest LMXB UV source (Willis, et al., 1978). A prototype of the LMXBs, Sco~X-1 consists of a neutron star and a low mass $(\\le 1 M_{\\odot})$ secondary in a close ($\\sim 10^{11}$ cm) 0.787d binary orbit (Crampton {\\it et al.} 1976). Sco~X-1 is one of two ``Z-class'' sources (displaying horizontal, normal, and flaring branches in the X-ray spectrum; see, e.g. Lewin and van der Klis, 1994, for a review) with an identifiable optical counterpart. The secondary (optical) star has a mass approximately 1 M$_\\odot$; the spectrum shows emission lines due to H and He II and no recognizable absorption features; if the mass transfer is driven by stellar evolution, then the star is at or near end of its main sequence age (Cowley, 1976; Cowley and Crampton, 1975; Gottlieb, Wright, and Liller, 1975). The increase in QPO frequency as the source traverses the `Z' shaped path in the X-ray color-color diagram has been intrepreted as being due to an increase in mass accretion rate. If so, the `normal branch' QPOs are characterized by an inverse behavior between mass accretion rate and X-ray luminosity. More recently, Sco X-1 has been shown to be one of the sources of high frequency (kHz) QPOs (van der Klis et al., 1996; 1997). The near-absence of stellar absorption features has slowed the interpretation of the optical and UV spectra of LMXBs. The mass transfer rates necessary to fuel the X-ray emission (assuming 10$\\%$ efficiency), together with the stellar masses and orbital periods, suggest a similarity between LMXBs and high mass accretion rate cataclysmic variables (i.e. nova-like variables; Cordova and Mason, 1982). However, the optical and UV luminosities of LMXBs exceed those of CVs by factors 10-100. This suggests that most of the optical and UV luminosity ($L_{opt} \\sim L_{UV} \\sim 10^{-2} L_X$) comes from the accretion flow as it is illuminated by the X-rays from the neutron star. The UV observations of Sco X-1 taken prior to September 1988 by the International Ultraviolet Observer (IUE) satellite have been summarized by us in an earlier paper (Kallman, Vrtilek, and Raymond, 1990, hereafter Paper I), and more recent IUE data by Vrtilek, et al. (1991). These data have been tested for variability associated with orbital motion, and for correlations among the various observables. The continuum spectra were found to fit to simple accretion disk atmosphere models which demonstrate that X-ray heating dominates the outer regions of the accretion disk. Correlations were found among the various emission line strengths, and upper limits to the emission line widths provided constraints on the location of the line emitting material if the emission is assumed to come from a Keplerian disk. IUE has shown that the UV emission line spectrum of Sco~X-1 (Paper I; Vrtilek {\\it et al.} 1991) is dominated by strong emission lines of C IV near 1550$\\AA$ and N V near 1238$\\AA$ and shows weaker lines of Si IV, O IV, O V, N IV, C III and less ionized species. The UV line strengths, ratios, and profiles vary noticeably ($\\sim 20\\%$) on the shortest timescales ($\\sim$0.5 hr) that can be probed with IUE; but IUE does not have the time resolution to detect variability in the UV at frequencies higher than those corresponding to this timescale. If such frequencies were accessible observationally we could obtain information about the behavior of the accretion flow within the Roche lobe of the primary. In spite of the effort devoted to understanding the X-ray variability behavior of LMXB's, little is known about their circumstellar environment and evolution. For example, estimates of the distance and reddening to Sco X-1 vary by a factor of 3 (Schachter et al., 1987; Willis, et al., 1980); secure distance estimates are vital to the understanding of the energetics of the X-ray emission. In addition, the X-rays and possible X-ray induced outflows from LMXB's are likely to affect the circumstellar environment by producing significant column densities of highly ionized species. Absorption studies, using Sco X-1 as a UV light source, can test for these ions. HST spectral resolution $\\Delta V \\sim 3 {\\rm km s}^{-1}$ (HRS Echelle A), enables the dissection of the circumstellar environment that X-ray, optical or IUE observations cannot do. In this paper we describe the results of observations of Sco X-1 in the ultraviolet and X-ray energy bands carried out in February 1996 using the HST and XTE satellites. The goals of these observations are: (i) Search for UV variability on timescales shorter than those previously accessible. The X-ray variability behavior has been extensively studied, and the X-ray power spectrum is known to extend to $\\geq$ 10 Hz during all of the variability states (e.g. Hasinger and Kurster, 1990). According to the reprocessing hypothesis the UV will reflect the variability of the X-rays on some timescales. For example, the accretion disk of Sco X-1 has a radius of about 2 lt-seconds so we expect that UV variability will be smeared on timescales much shorter than this; the UV emitting region should act as a ``low pass filter'' to the X-ray power spectrum. (ii) With HST, we can measure the interstellar column density toward Sco X-1 using interstellar L${\\alpha}$, thereby providing a third, independent distance estimate, which may help resolve the uncertainty. The interstellar line is broader than either the intrinsic L$\\alpha$ emission from Sco~X-1 or the geocoronal line, and is apparent in the IUE spectra. (iii) Sco X-1 is likely to have an effect on any nearby interstellar medium. This can occur by X-ray photoionization, which will produce significant column densities of ions such as C IV and N V within a radius of $\\sim$3 pc (McCray, Wright, and Hatchett, 1976). These will have a detectable signature in high resolution observations as narrow features, with widths of $\\sim 0.1 \\AA$. The relative strengths of such features provide a measure of the ionization balance in the interstellar medium, which in turn can constrain the density and X-ray flux history. In what follows we describe our attempt to test some of these ideas using simultaneous observations of Sco X-1 by HST and XTE. In section 2 we describe the observations themselves; in sections 3 and 4 we discuss the spectral and timing analysis, respectively. In section 5 we discuss the results and summarize the conclusions. \\vfill \\eject ", + "conclusions": "" + }, + "9802/astro-ph9802064_arXiv.txt": { + "abstract": "The quality, depth, and multi--color nature of the Hubble Deep Field images makes them an excellent resource for studying galaxies at $z > 2$ using selection techniques based on the presence of the 912\\AA\\ Lyman break. I present a descriptive review of this method and of the properties of the objects which it identifies, and summarize spectroscopic progress on galaxies with $2 < z < 4$ in the HDF. Using ground--based and HDF samples of Lyman break galaxies I discuss the luminosity function of galaxies at $z \\approx 3$, and consider the effects of extinction on the star formation rates that are derived from the UV luminosity information. Infrared observations of the HDF provide data on the rest--frame optical properties of $z \\approx 3$ galaxies, which are briefly described. ", + "introduction": "Although the study of galaxies at high redshift neither begins nor ends with the Hubble Deep Field (HDF), this conference has demonstrated the ways in which the HDF has served to focus the attention of the community on the properties of galaxies at $z > 2$. In part this is because the HDF imaging data was obtained through several filters, permitting the use of color selection techniques to isolate and study populations of galaxies at various redshifts. Because the HDF images are so deep, colors can be measured with unusually high precision (and in a spatially resolved fashion {\\it within} individual galaxies) for objects which ordinarily would be considered very faint for ground--based telescopes. Additionally, the HDF images easily detect galaxies at magnitudes well beyond the spectroscopic limits of even the largest telescopes. The desire to understand their nature requires {\\it some} idea of their distances, and it is therefore tempting to look for photometric means of estimating redshifts without the benefit of ordinary spectroscopy. For these reasons, many presentations at this symposium have considered various applications of ``photometric redshift'' estimation in the HDF and in other data sets. One such method takes advantage of the ubiquitous 912\\AA\\ Lyman limit discontinuity, which is redshifted into the HDF bandpasses at $z \\simgt 2$. In recent years, color selection based on the Lyman limit has developed into a highly successful means of detecting galaxies at large redshifts, as I will review below. In designing the HDF observations, our working group at STScI incorporated F300W imaging into the four--filter scheme for two reasons, one scientific and one purely practical. First, such data offers the potential for Lyman break selection of high redshift galaxies. Second, we wished to take advantage of the ``bright'' portions of the orbit during Continuous View Zone visibility, when scattered earthlight severely impacts WFPC2 imaging at redder wavelengths. The reduced amplitude of the scattered background at UV wavelengths, and the fact that F300W imaging with WFPC2 is not normally background limited anyway, made bright--time observing through that bandpass a suitable use for these otherwise disadvantageous observing intervals. I begin with a descriptive review of the Lyman break technique, illustrated using data from the HDF. Although the HDF Lyman break galaxy sample is much smaller than that which has been identified from ground--based imaging studies (which cover much larger solid angles), the excellent photometric precision of the HDF data and the large percentage of Lyman break galaxies present at the fainter magnitude limits which it probes makes it quite useful for illustrating the principles of the method. In \\S 2.1 I summarize the current status of spectroscopy on HDF Lyman break galaxies. In \\S 3 I discuss some of the statistical properties of Lyman break galaxies, concentrating on their numbers, luminosities, and colors, and inferences that can be derived from these measurements. Finally, \\S 4 discusses the rest--frame optical properties of $z > 2$ galaxies in the HDF using deep infrared imaging data. In his contribution to this volume, Mauro Giavalisco discusses our ground--based Lyman break galaxy sample in greater detail, and addresses the spectroscopic, morphological, and clustering properties of Lyman break galaxies, which I will largely neglect here. In this paper, HDF galaxy selection and object names are based on the catalog of Williams et al.\\ (1996). Optical photometry is reported on the AB magnitude system, with the WFPC2 bandpasses indicated by \\U300, \\B450, \\V606 and \\I814. In general I have used new, revised photometry of the Williams et al.\\ catalog objects based on optimized apertures which maximize signal--to--noise for color measurement. This has the advantage of allowing robust Lyman break galaxy selection to somewhat fainter magnitudes than was previously possible. ", + "conclusions": "In this presentation I have tried to use HDF data to illustrate the ways in which multi--color photometry can be used first to select galaxies at high redshift and then to learn something about their intrinsic properties. The great depth of the HDF and the precision of its photometry makes it ideal for illustrating and applying the Lyman break color selection technique: cf.\\ the remarkable prominence of the high redshift ``plume'' in figures 2 and 4. By \\V606 = 26.5, nearly 1 in 4 HDF galaxies is a Lyman break candidate and thus is likely to be at $z > 2$. Spectroscopy of HDF galaxies at all redshifts has proceeded at a remarkable pace. After only two observing seasons, the central HDF is almost certainly the piece of celestial real estate with the highest surface density of measured galaxy redshifts (now $\\approx 24$/arcmin$^{2}$, of which $\\sim$20\\% are at $z > 2$). The distribution of ultraviolet luminosities of $z \\approx 3$ galaxies, converted to star formation rates using a simple prescription, spans a similar range to that of galaxies in the local universe. Schechter function fits give a characteristic 1500\\AA\\ specific luminosity of $M_{AB} \\approx -21$ (for $H_0 = 50$, $q_0 = 0.5$), corresponding to a star formation rate of $\\sim 14 M_\\odot$/year. Very few Lyman break galaxies have ``raw'' SFRs $> 50 M_\\odot$/year. If these measurements are taken at face value, these objects cannot produce the total stellar mass of $\\sim L^\\ast$ galaxies in short timescales as traditional ``monolithic'' formation scenarios for elliptical galaxies would require. This, then, would suggest that massive galaxy formation takes place either at still higher redshifts where we have yet to look, or proceeds by the hierarchical assembly of smaller objects as expected in theories such as CDM. However, the possible extinction corrections to the UV luminosities of Lyman break galaxies are highly uncertain and could be quite large. Their UV colors are redder than expected from spectral models of ``naked'' star forming galaxies, a fact which could easily be explained by the presence of dust. The derived extinction corrections based on these colors, however, are extremely sensitive to the form of the extinction law in the ultraviolet, and can range from factors of 2 to $> 7$. New observations in the infrared, and eventually at far--IR and sub--millimeter wavelengths, will provide an independent measure of star formation rates which can be useful for addressing this question. Characteristic rest--frame optical luminosities of Lyman break galaxies, as measured from infrared photometry, are $M_V = -22$. Their UV--optical rest--frame colors galaxies span a range which would be typical for normal spirals in the nearby universe. Detailed morphological study of $z \\approx 3$ galaxies at rest--frame optical wavelengths must await observations with NICMOS, which will take place in 1997--1998. Although Lyman break color selection is a simple technique compared to more sophisticated photometric redshift methods, it has the virtue of being relatively model independent and easy to apply and understand. Like all such methods, however, its utility depends strongly on the degree to which it is tested and calibrated by follow--up spectroscopy, which requires substantial effort on large ground--based telescopes. With hundreds of Lyman break redshifts in hand, we can begin to carry out fairly sophisticated analyses of the luminosity function, clustering, and other properties of galaxies at $z \\approx 3$. For the HDF, despite the impressive observing efforts to date, we are unlikely ever to succeed in collecting hundreds of redshifts for $z > 2$ galaxies with existing telescopes and instrumentation. However, we may take advantage of the insights gained from studying the Lyman break galaxy population in non--HDF data sets to aid interpretation of the HDF objects, and thus to use the HDF to push the method to different flux and redshift limits. Regardless of how much we learn about high redshift galaxies in the HDF, we must remind ourselves of what a small volume the HDF probes. The entire comoving volume over which \\U300--dropouts have been found in the HDF, $2 < z < 3.5$, is only 18000$h_{50}^{-3}$~Mpc$^3$ for $q_0 = 0.5$. Locally, that would correspond to a sphere with radius 16.2$h_{50}^{-1}$ Mpc -- not even reaching the Virgo cluster! We must therefore be cautious about how representative HDF galaxies are in any statistical sense, particularly in light of recent evidence for strong clustering at $z > 2$. The pencil--beam geometry of the HDF volume ensures that it will traverse a wider range of large scale structures than would the corresponding 14.2~Mpc radius sphere locally, but such a geometry brings its own complications for some applications. For example, clustering may introduce large fluctuations in $N(z)$ which can seriously complicate analyses of the redshift evolution of galaxy properties, global luminosity density, etc. It is therefore risky to extrapolate too far from the HDF to the properties of galaxies in the high redshift universe as a whole. Ultimately, however, the insights gathered from the HDF, when calibrated with data from ground--based, large--volume surveys, should provide a powerful means of understanding the early stages in the evolution of normal galaxies." + }, + "9802/astro-ph9802252_arXiv.txt": { + "abstract": "We analyze the monopole in the peculiar velocities of 44 Type Ia supernovae (SNe Ia) to test for a local void. The sample extends from $20$ to $300\\hmpc$, with distances, deduced from light-curve shapes, accurate to $\\sim 6\\%$. Assuming $\\omm=1$ and $\\oml=0$, the most significant deviation we find from the Hubble law is an outwards flow of $(6.6\\pm2.2)\\%$ inside a sphere of radius $70\\hmpc$ as would be produced by a void of $\\sim 20\\%$ underdensity surrounded by a dense shell. This shell roughly coincides with the local Great Walls. Monte Carlo analyses, using Gaussian errors or bootstrap resampling, show the probability for chance occurrence of this result out of a pure Hubble flow to be $\\sim 2\\%$. The monopole could be contaminated by higher moments of the velocity field, especially a quadrupole, which are not properly probed by the current limited sky coverage. The void would be less significant if $\\omm$ is low and $\\oml$ is high. It would be more significant if one outlier is removed from the sample, or if the size of the void is constrained \\apriori. This putative void is not in significant conflict with any of the standard cosmological scenarios. It suggests that the Hubble constant as determined within $70\\hmpc$ could be overestimated by $\\sim 6\\%$ and the local value of $\\Omega$ may be underestimated by $\\sim 20\\%$. While the present evidence for a local void is marginal in this data set, the analysis shows that the accumulation of SN Ia distances will soon provide useful constraints on elusive and important aspects of regional cosmic dynamics. ", + "introduction": "\\label{sec:intro} Large-scale redshift surveys of galaxies show underdense regions of typical extent $\\sim 50\\hmpc$. These ``voids\" appear to be bordered by dense ``walls\" (Kirshner \\etal 1981; Huchra \\etal 1983 [CfA]; Broadhurst \\etal 1990; Shectman \\etal 1996 [LCRS]). In particular, maps of our cosmological neighborhood display the Great Wall of Coma and the Southern Wall, that appear to connect into a shell-like structure of radius $70-80\\hmpc$ about the Local Group (Geller \\& Huchra 1989 [CfA2]; da Costa \\etal 1994 [SSRS2]). The volume encompassed by this structure appears to be of lower density. Despite these impressive maps, it is difficult to quantify the large-scale radial density profile of this region. First, the true galaxy density is hard to distinguish from the sample selection function when the structure of interest approaches the sample size. Second, we do not know how well galaxies trace mass. And third, portions of the galaxy distribution are obscured from our viewpoint within the Milky Way. The imprint of wall and structure on the power spectrum may possibly be associated with excess power observed at a wavelength of $\\sim 100-150\\hmpc$ (\\eg, Broadhurst \\etal 1990; Landy \\etal 1996; Einasto \\etal 1997). This scale might be naturally attributed to the scale of the cosmological horizon after the universe became matter dominated and before the plasma recombined, though this peak is only expected to be significant for relatively high values of the baryon density (Hu \\& White 1997). A local void has been proposed as one way to reconcile the age of the Universe based on the Hubble expansion with the ages of globular clusters within the framework of the Einstein-de Sitter cosmology (\\eg, Turner, Cen \\& Ostriker 1992; Bartlett \\etal 1995) Measurements of the Hubble constant within the void would overestimate the universal value by $\\delta \\rho/\\rho \\approx -3\\, \\delta H/H$. Indeed, the values obtained for the Hubble constant from the longest-range distance indicators, the SNe Ia (Jacoby et al 1992; Sandage \\& Tammann 1993; Tammann \\& Sandage 1995; Hamuy et al 1995, 1996b; Riess, Press, \\& Kirshner 1995a, 1996; Branch, Nugent, \\& Fisher 1996) and the gravitational lenses (Falco \\etal 1997; Keeton \\& Kochanek 1997) are typically smaller than values obtained more locally using Tully-Fisher (TF) distance indicators (Kennicutt, Freedman, \\& Mould 1995; Mould \\etal 1995; Freedman \\etal 1994; Freedman 1997). A local void would also imply that local estimates of $\\Omega$ underestimate the global value of $\\Omega$. Finally, a local outflow would reduce the distances derived from TF peculiar velocities for features such as the Great Attractor, bringing them into better agreement with the positions derived from redshift surveys (Sigad \\etal 1998). It is important to separate impressions and theoretical wishes from firmly established observational facts. Attempts to establish monopole deviations from pure Hubble flow have not yet produced conclusive results. For example, Shi (1997) claimed finding a monopole gradient in subsamples of the Mark III catalog of Tully-Fisher peculiar velocities (Willick \\etal 1997a). But he also found a marginal monopole gradient of the opposite sign in the peculiar motions of rich clusters based on their brightest members as standard candles (Lauer \\& Postman 1992; 1994). Finally, he found no significant deviation in an early subsample of 20 SNe (Riess \\etal 1996). None of these data sets was ideal for testing for a void of radius $\\sim 70\\hmpc$: the Mark III data include only a small number of galaxies beyond $70\\hmpc$, the earlier SNe sample is too sparse, and the rich clusters, while of comparable abundance to the current SNe sample within $70\\hmpc$, have a much larger error per cluster. Kim \\etal (1997), based on comparing SNe distances at high redshifts ($z\\sim 0.4$) and low redshifts ($z < 0.1$), put a $2\\sigma$ upper limit of $10\\%$ for an outflow within the local $300\\hmpc$, but their sample has only a few SNe within $70\\hmpc$, which is the suggested domain of the putative void. We now have a sample of 44 SNe Ia from the Calan/Tololo Survey (Hamuy et al 1993, 1996a) and from the CfA supernova program (Riess 1996, Riess \\etal 1998), which reaches the threshold for an interesting assessment of the dynamical signature of a void contained within the great walls. The nearest SNe in this sample are at distances of $\\sim 20\\hmpc$, the furthest are beyond $300\\hmpc$, the median is at $\\simeq 95\\hmpc$, and 17 SNe lie inside $70\\hmpc$. This distribution in distance is well suited for detecting a monopole perturbation within $50-100\\hmpc$ while determining the universal Hubble flow outside this sphere. These are the scales which have been suggested and where the standard cosmological theories predict perturbations on the order of several percent (\\eg, Wang, Spergel \\& Turner 1998). Because SNe light curves provide distances with a precision of $5-8\\%$, we believe it is worth carrying out a preliminary analysis with only 44 tracers. This sample has already been used to put constraints on the global Hubble flow and the dipole motion (Hamuy et al 1995, 1996b; Riess, Press, \\& Kirshner 1995a,b, 1996), and the nearby sample has been used to assess how well the velocities of SNe match the gravity field inferred from galaxy redshift surveys (Riess \\etal 1997). In \\S~\\ref{sec:data} we describe the data. In \\S~\\ref{sec:monopole} we measure the peculiar monopole and evaluate its significance in an Einstein-de Sitter Universe. In \\S~\\ref{sec:cautions} we explore the effect of different cosmologies, error estimates and higher velocity-field moments on our results. In \\S~\\ref{sec:discussion} we discuss the implications of our results. ", + "conclusions": "\\label{sec:discussion} We have analyzed the monopole perturbations in the peculiar velocity field as probed by 44 SNe Ia with small distance errors. For an Einstein-de Sitter universe, we find an outward perturbation in the Hubble flow of $(6.5 \\pm 2.2)\\%$ within a sphere of radius $70 \\hmpc$, with a possible indication for inflow in the surrounding shell. The void is less significant if $\\omm$ is low and $\\oml$ is high. Its significance is increased when one outlier is excluded, or when the size of the void is constrained \\apriori to a limited range. In the idealized case of a spherically symmetric density distribution, an observed monopole in the velocity field is associated with a specific density profile around the Local Group. In linear theory, $\\delta \\rho/\\rho = -\\nabla\\cdot\\vv/\\Omega^{0.6}$, so that a constant velocity monopole $\\delta_H$ corresponds, for $\\Omega=1$, to a constant density fluctuation offset of $\\delta \\rho/\\rho = - 3\\, \\delta_H$. In this picture, our results would require a local void of $\\sim -20\\%$ density contrast. The sharp drop in $\\delta_H$ at $R=70-75\\hmpc$ would be produced by a narrow shell of high mass density at that distance. This is approximately the observed location of the shell-like structure defined in redshift surveys by the Great Wall of Coma and the Southern Wall. It is tempting to associate the drop in the outflow we see with these features of the regional galaxy distribution. Indeed, Dell'antonio, Geller \\& Huchra (1996), based on a sample of peculiar velocities of spiral galaxies in the vicinity of the Great Wall, estimate the real-space width of the wall to be smaller than $11\\hmpc$ at $90\\%$ confidence. Their corresponding limit on the overdensity of the Great Wall is $\\delta \\rho/\\rho < 2/ \\Omega^{0.6}$. An outer region of negative mean $\\delta_H$ beyond $75\\hmpc$ would represent back-streaming from voids behind the walls into the walls. We expect that the density profile averaged over a sphere returns to its mean value ($\\delta M/M=0$) where the velocity monopole converges to the global Hubble flow. The evidence for such a local void in density profiles derived from redshift surveys is not conclusive. While we see no evidence for a local void in the IRAS 1.2 Jy survey as analysed by Koranyi \\& Strauss (1997, Fig.\\ 6), Fig.\\ 5 of Springel \\& White (1997) clearly shows evidence for a void based on the same data. A similar void is seen in Saunders \\etal (1990, Fig.\\ 11) based on the QDOT-IRAS redshift survey. Preliminary results from the PCSZ-IRAS redshift survey provide further evidence for the presence of a local void of even larger extent (S.D.M.\\ White 1997, private communication). The apparently conflicting results could be partly explained by the fact that the effective survey size, especially in the 1.2 Jy survey, is not much larger then the volume of the local void. This might introduce some confusion between the radial selection function and the real density profile. Because of the small number of data points in the current sample, we have not attempted to determine the center point about which the angle-averaged void is deepest or largest. This center point does not necessarily coincide with the Local Group. Such a void is not in significant conflict with expectations from standard cosmological theories. The expected \\rms perturbations on these scales, as predicted by popular CDM models, range from $2\\%$ to $4\\%$, with standard-CDM ($\\omm=1$) being on the high side and Open CDM ($\\omm=0.3$, $\\oml=0$) on the low side (see also Shi \\& Turner 1998). The local void we find, of $\\sim 6 \\pm 2\\%$, is thus less than a $2 \\sigma$ deviation from the predictions of any of the conventional models. Although an open model is less likely than the others, the void in this case is slightly shallower and corresponds only to a $\\sim 1.9\\sigma$ deviation from the predictions of an open CDM model. One should also worry about possible observational systematics that may be artificially interpreted as a local peculiar outflow. Sources of observational error that vary systematically with distance are prime suspects. These include the effects of redshift on the integrated spectral light which passes through a fixed broad band filter, the ``K-correction''. Although care has been taken to include such corrections, the non-stellar features of a SN Ia spectrum inhibit an exact compensation for this effect. However, the magnitude of the correction is fairly small for the bulk of our SNe data, typically less than $0.02$ mag (Hamuy \\etal 1993b). This corresponds to a peculiar velocity of $\\sim 50\\kms$, which is less than $1\\%$ of deviation from Hubble flow within the local void compared to the detected signal of $\\sim 6\\%$ deviation. Our estimate of $10\\%$ error in the K-correction implies that these effects are negligible. Similarly, the range of supernova luminosities observed in the nearby sample is larger than at great distances because it includes intrinsically faint objects. Any defects in the ability of MLCS to correctly account for these systematic shifts could, in principle, cause problems in making reliable inferences about the void dynamics. However, this does not seem to be an important effect. The faint tail in the inner region is populated by only three SNe, of which only one, the brightest of the three, has a significant positive $\\delta H / H$. We have also tested for correlations between the distances or peculiar velocities and certain quantities that might have systematically affected them These quantities include the SN detection time with respect to the time of the light-curve maximum, the absolute correction to the light-curve template, and the magnitude of the extinction correction. We have found no significant correlations. One implication of our result is that the local value of the Hubble constant cannot be {\\it much} larger than the global value. Our best estimate of $\\delta_H$ is $(6.5 \\pm 2.2)\\%$, so if the regional value (inside $70\\hmpc$) were $70\\kmsmpc$, the global value would be $65\\kmsmpc$. But for the global value to reach $50\\kmsmpc$ would require $\\delta_H$ of $40\\%$, which is clearly excluded by our analysis. This assumes that the volume sampled by our set of SNe Ia reaches out to the mean density and that we are not embedded in a deeper void of even larger extent. A possible hint for such a void may come from the observed over-abundance in faint galaxy counts, which may, or may not, be explained by luminosity evolution of galaxies (\\eg, Heydon-Dumbelton, Collins, \\& McGillivray 1989; Maddox et al 1990; Lilly 1993; Driver, Windhorst, \\& Griffiths 1995; Gronwall \\& Koo 1995). One controversial interpretation is of a significant deficiency of galaxies on size scales of $\\sim 300\\hmpc$ (Huang et al 1997). However, in the Las Campanas Redshift Survey, the galaxy density profile clearly converges to a universal value beyond $150\\hmpc$ (Lin \\etal 1997, Fig. 8). Furthermore, a comparison of supernovae distances at low and high redshifts puts an upper limit of $10\\%$ at $95\\%$ confidence on the depth of a void of size $\\sim 300\\hmpc$ (Kim \\etal 1997). Our tentative local void detection is also consistent with upper limits on $\\delta_H$ of $\\sim 7\\%$ obtained from Abell/ACO clusters (Lauer \\& Postman 1992; Lauer \\etal 1997). The $\\sim 6.5\\%$ outflow detected here would partially help reconcile the large-scale estimates of the Hubble constant, \\eg, from SNe (Tammann \\& Sandage 1995; Riess \\etal 1996) and estimates that are limited more to the inner volume, such as based on Tully-Fisher distances (Kennicutt, Freedman, \\& Mould 1995, Mould \\etal 1995, Freedman \\etal 1994, Freedman 1997). The latter are typically $\\sim 10-20\\%$ higher. The local void marginally detected here would help, in a humble way, in the resolution of the apparent conflict between the Hubble constant and the ages of globular clusters (Bolte \\& Hogan 1995; Chaboyer \\etal 1998). The implication for the value of $\\Omega$ is that estimates based on data within the local void (see a review by Dekel, Burstein, \\& White 1997) would underestimate the universal value by $\\sim 20\\%$. We conclude by reiterating that our detection of the local void is only marginal. Since its statistical significance can be interpreted in several different ways, we presented a ``supermarket\" of statistical evaluations. If we were investigating the properties of a void whose scale had been established by other data, such as the great walls seen in redshift slices and pencil beams, we would report a $\\sim 6.5\\%$ outflow which, in an Einstein de-Sitter universe, has a probability of order $1\\%$ to be obtained by chance from a pure Hubble flow. The void would be even more significant if one outlier were removed. However, with no pre-conception about the extent of the void, the detection is less decisive, with the probability for a chance occurrence rising to a few percent. For the low $\\omm$ and high $\\oml$ models this probability is as high as $10\\%$, and in every case there is possible confusion with higher moments of the velocity field, especially a quadrupole. In any case, our indication of a modest local void is tentative and should be confirmed (or refuted) by future observations; our detection poses a specific model to be tested, that of an outflowing region of radius $70\\hmpc$." + }, + "9802/astro-ph9802314_arXiv.txt": { + "abstract": " ", + "introduction": "The phenomenon of counterrotation is observed when two galaxy components have their angular momenta projected antiparallel onto the sky. It follows that if the two components rotate around the same axis, the counterrotation is intrinsic. On the contrary the counterrotation is only apparent if the rotation axes are misaligned and the line-of-sight lies in between the two vectors or their antivectors. In the case of intrinsic counterrotation the two components can be superimposed or radially separated. As far as the two components are concerned, stars are observed to counterrotate with respect to other stars or gas. The counterrotation of gas versus gas has been also detected. Up to now, the number of galaxies exhibiting these phenomena are $\\sim60$, the morphological type of which ranges from ellipticals to S0's and to spirals. Previous reviews about counterrotation are those of Rubin (1994b) and Galletta (1996). When a second event occurs in a galaxy, such as the acquisition of material from outside, it is likely that the resulting angular momentum of the acquired material is decoupled from the angular momentum of the preexisting galaxy. Counterrotation is therefore a general signature of material acquired from outside the main confines of the galaxy. Good examples of such cases are ellipticals with a dust lane or gaseous disk along the minor axis and polar ring galaxies, where the angular momenta are perpendicular. It should be noted that recently attempts have been made to explain special cases of stars versus stars counterrotation in disk galaxies as due to a self induced phenomenon in non-axisymmetric potentials (Evans \\& Collett 1994, Wozniak \\& Pfenninger 1997). In the following we discuss the phenomenon according to the morphological type and to the kind of counterrotation. ", + "conclusions": "Counterrotation occurs in a wide variety of forms (gas vs. stars, stars vs. stars, gas vs. gas). It is present in galaxies of different morphological types, ranging from ellipticals to early-type disk galaxies. The counterrotation of stellar vs. stellar disks is the type of counterrotation we expect to be prevailing, since it is the end result of gas vs. stars counterrotation. Therefore its frequency could be very high if the area described in Fig.~1 is carefully inspected, using the state of the art analysis of the shape and asymmetries of the line profiles of the absorption lines. This would indicate that acquisition and merging events are common phenomena in the history of galaxies.% We are grateful to D. Burstein for useful comments on the manuscript." + }, + "9802/astro-ph9802136_arXiv.txt": { + "abstract": "A novel scheme for measuring the cross section of the $^7$Be(p,$\\gamma)^8$B reaction, the major source of high energy neutrinos from the sun, is presented. The scheme involves a strictly uniform particle beam and overcomes some of the recognized experimental uncertainties of previous measurements. A new measurement of $\\sigma[^7$Li(d,p)$^8$Li] has been carried out using this setup, and the present value of $\\sigma[^7$Li(d,p)$^8$Li] = 155(8) mbarn at the top of the E$_d$(lab.) = 776 keV resonance is compared to previous measurements. A new issue regarding both the (d,p) and (p,$\\gamma$) reactions has been examined: reaction-product nuclei which are backscattered out of the target. Measurements and simulations carried out in the course of this investigation are presented and discussed in the context of possible effects on the measured cross sections of these reactions. PACS Numbers: 95.30.Cq; 96.60.Kx; 25.45.Hi ", + "introduction": "\\par The cross section of the reaction $^7$Be(p,$\\gamma)^8$B has recently become the subject of renewed intense scrutiny [1] owing to the pre-eminence of the $^8$B reaction products as the main source of high energy neutrinos in the interior of the sun and, hence, the possible implications for the so-called ``solar neutrino problem\". A number of new precision measurements of this cross section in various stages of preparation have been reported lately. We have set up an experiment for such a measurement at the Weizmann Institute, focusing on one of the major sources of uncertainty in previous experiments: the homogeneity of the areal density of the target material. \\par In general, when a nonuniform particle beam impinges on a nonuniform target, the reaction yield is given by: \\be Y=\\sigma\\int{dn_b\\over dS}{dn_t\\over dS}dS \\ee where $n_b,n_t$ are the respective total numbers of beam and target particles and ${dn_b\\over dS},{dn_t\\over dS}$ are areal densities. \\\\ Only when the target is known to be uniform and the beam is smaller than the target can eq.(1) be simplified to: \\be Y=\\sigma {dn_t\\over dS}\\int{dn_b\\over dS}dS = \\sigma {dn_t\\over dS}n_b \\ee In such a case, the evaluation of the cross section is independent of the areal distribution of the particle beam. On the other hand, in other cases - e.g. for radiochemically produced $^7$Be targets [2] - the target cannot be assumed to be strictly uniform and the full relation (1) has to be used in the evaluation. The inherent uncertainties in the distributions ${dn_b\\over dS}$ and ${dn_t\\over dS}$ may thus lead to considerable uncertainties in the value of the integral and hence in the deduced cross section. We have addressed this problem by inverting the arrangement: we use a homogeneous beam - produced by raster scanning - impinging on a target {\\underline {smaller}} than the beam. The relation (1) then reduces to: \\be Y=\\sigma {dn_b\\over dS}n_t \\ee independent of the potentially problematic determination of the target areal distribution. \\par As a first step we measured the cross section of the reaction $^7$Li(d,p)$^8$Li at the top of the resonance at E$_d$(lab.)=776 keV. The mirror nuclei $^8$Li and $^8$B have very similar mean lives and similar $\\beta$ decays ($\\beta^-$ and $\\beta^+$, respectively) to $^8$Be $\\rightarrow$ 2$\\alpha$. The main differences are (i) the much easier preparation of a $^7$Li target and (ii) the much higher yield from $^7$Li(d,p). The reaction $^7$Li(d,p)$^8$Li can therefore provide a convenient check of the equipment and the method. Beyond that, this reaction is of significance in its own right and as a stepping stone and calibration in some $^7$Be(p,$\\gamma$)$^8$B experiments [1-4].\\\\ \\par In the course of the present investigation we have examined a new issue: reaction-product nuclei being backscattered from the target backing, leaving the target assembly and reducing the number of detected $\\alpha$ particles. We present below experimental results which probe this issue and compare them to computer simulations. The influence of this effect on past and future measurements is discussed. ", + "conclusions": "" + }, + "9802/astro-ph9802246_arXiv.txt": { + "abstract": "Absorption due to He~II \\Lya\\ has now been detected at low resolution in the spectra of four quasars between redshifts $z =$ 2.74 -- 3.29. We assess these observations, giving particular attention to the radiative transfer of the ionizing background radiation, cloud diffuse emission and ionization physics, and statistical fluctuations. We use high-resolution observations of H~I absorption towards quasars to derive an improved model for the opacity of the intergalactic medium (IGM) from the distribution of absorbing clouds in column density and redshift. We use these models to calculate the H~I and He~II photoionization history, the ratio $\\eta =$ He~II/H~I in both optically-thin and self-shielded clouds, and the average line-blanketing contribution of the clouds to He~II absorption. The derived ionization rate, $\\Gamma_{HI} = (1-3) \\times 10^{-12}$ s$^{-1}$ ($z = 2-4$) is consistent with the ionizing background intensity inferred from the ``proximity effect'', but it remains larger than that inferred by N-body hydrodynamical simulations of the Ly$\\alpha$ absorber distribution. The He~II observations are consistent with line blanketing from clouds having $N\\subH \\geq 10^{12}$ cm$^{-2}$, although a contribution from a more diffuse IGM would help to explain the observed opacity. We compute the expected He~II optical depth, $\\tauhe(z)$, and examine the implications of the sizable fluctuations that arise from variations in the cloud numbers and ionizing radiation field. We assess how He~II absorption constrains the intensity and spectrum of the ionizing radiation and the fractional contributions of the dominant sources (quasars and starburst galaxies). Finally, we demonstrate how high-resolution ultraviolet observations can distinguish between absorption from the diffuse IGM and the Ly$\\alpha$ forest clouds and determine the source of the ionizing background. ", + "introduction": "Absorption in the He~II $\\lambda$303.78 \\Lya\\ line has long been considered a potentially important tool for studying the high-redshift universe (\\ME\\ \\& Ostriker 1990; \\ME\\ 1993; Shapiro, Giroux, \\& Babul 1994). As early estimates showed, He~II is more abundant than H~I in the highly ionized ``\\Lya\\ forest clouds'' that appear in quasar spectra, because He~II has a lower photoionization cross section and a larger recombination rate. Comparison of the H~I and He~II absorption could tell us about the photoionizing background, the history of structure formation, and internal conditions in the \\Lya\\ clouds. Also, He~II absorption from a smoothly distributed intergalactic medium (IGM) should be more easily detectable than the corresponding H~I absorption, which has proven difficult to measure unambiguously (Giallongo \\ea 1994; Williger \\ea 1994; Fang \\& Crotts 1995). These hopes have partially been realized by recent observations of He~II toward four quasars. \\Qj\\ was observed with two instruments on the Hubble Space Telescope (HST): at low resolution with the Faint Object Camera (FOC) (Jakobsen \\ea 1994) and at higher resolution with the Goddard High Resolution Spectrograph (GHRS) (Hogan, Anderson, \\& Rugers 1997). \\Qt\\ was observed with HST using both the Faint Object Spectrograph (FOS; Tytler \\ea 1995) and the FOC (Tytler 1997). \\HS\\ was observed with the Hopkins Ultraviolet Spectrometer (HUT) during the Astro-2 mission (Davidsen \\ea 1996, hereafter D96). Most recently, \\HE\\ was observed with HST/GHRS (Reimers \\ea 1997). The resolution of the HUT spectrum is $\\sim \\!3$~\\AA\\, while that of the HST/FOS spectra is $\\sim\\!15$~\\AA\\ in the FOC spectra. The HST/GHRS spectra potentially have higher resolution, although the data are usually binned to enhance the signal-to-noise ratio. Thus, individual absorption lines are not resolved. Instead, the results are reported (see also Figure~\\ref{fig:obs-opt-depth}) in terms of an average optical depth: $1.3 \\leq \\tauhe \\lesssim 4$ towards \\Qj\\ ($\\zem =3.29$), $1.0 \\leq \\tauhe \\leq 2.0$ ($\\zem = 3.18$) towards \\Qt, $\\tauhe = 1.00 \\pm 0.07$ towards \\HS\\ ($\\zem = 2.74$), and $\\tauhe$ varying from $\\sim\\!1$ to $\\geq4.8$ toward \\HE. As shown by the error estimates, the HUT result is the most restrictive of the four, and we normalize our plot to that point. We anticipate our discussion in \\S 3--4 by placing these data points in the context of two simple models for the evolving He~II Ly$\\alpha$ opacity. As we indicate in Fig. 1 and will discuss in detail below, a single high-precision measurement is inadequate to constrain models for the He~II opacity since the intrinsic fluctuations in the opacity at a given epoch are significant. Early interpretation of these results in the discovery papers, as well as by Madau \\& Meiksin (1994), Zheng \\& Davidsen (1994), Giroux, Fardal, \\&~Shull (1995, hereafter Paper I), and Songaila, Hu, \\&~Cowie (1995), has focused on the requirements for reproducing the observed He~II absorption. As yet, it has been difficult to place strong constraints on the required assumptions. These include the distribution of absorption-line clouds in redshift, H~I column, and H~I line width, and the populations of ionizing sources and their intrinsic spectra. Assumptions about the physical properties of the clouds must also be made, both to derive the ionization correction necessary to convert $\\nhone$ to $\\nhetwo$ and to estimate the distribution of He~II line widths. In addition, an unknown amount of He~II absorption may arise in diffuse gas in the IGM, i.e. the true Gunn-Peterson effect (Gunn \\& Peterson 1965, hereafter GP). These uncertainties also affect the analysis of H~I and He~II absorption by numerical hydrodynamical simulations (cf.\\ Zhang \\ea 1997). A further complication, discussed by Zheng \\& Davidsen (1995) and in Paper I, is that ionizing radiation from the observed quasar itself may alter the level of He~II ionization in nearby clouds. This ``He~II proximity effect'' is expected to be larger than the corresponding H~I proximity effect if, as is believed, intervening clouds strongly depress the level of He~II-ionizing radiation for the general metagalactic background. This effect is not obvious in the spectrum of \\HS, but it has been claimed in the spectrum of \\Qj\\ (Hogan \\ea 1997; Heap 1997). Fortunately, a substantial reduction in the volume of model parameter space is possible. A high-resolution \\Keck\\ spectrum of \\Qj\\ (Songaila, Hu, \\& Cowie 1995) has provided the actual set of H~I absorption lines to which the observed He~II line-blanketing absorption must correspond. It is likely that similar line lists will be generated for the other quasars used to measure He~II absorption. In addition, high-quality spectra along the lines of sight to many quasars are providing increasingly good statistics on the average distribution in column density, Doppler $b$-values, and redshift of H~I absorbing clouds. Additional He~II absorption observations are planned with HST, and significant advances are expected to come from the Far Ultraviolet Spectroscopic Explorer (FUSE), scheduled for launch in late 1998, and the Cosmic Origins Spectrograph (COS) scheduled for the HST refurbishment mission in late 2002. The FUSE spectrograph may be able to resolve individual lines at wavelengths down to 915 \\AA\\ ($z\\subHe \\geq 2.01$), providing line widths and directly giving the ionization correction in the absorbing clouds. The COS instrument offers significantly greater throughput, to study faint QSOs at $z\\subHe > 2.9$. The purpose of this paper is to ask whether any simple model fits all of the current observations and to identify important questions that could be answered in the future. Our approach is phenomenological; we do not employ a fundamental theory of the \\Lya\\ forest, as has been attempted in Gnedin \\& Hui (1996) or Bi \\& Davidsen (1997), nor do we use a numerical hydrodynamical approach. Instead, we employ the standard, observationally oriented division between ``clouds'' and a diffuse IGM. In \\S~2, we outline our calculation of the cosmological radiative transfer of direct and diffuse photoionizing radiation through the absorbing clouds. We compile a new absorption-line sample and construct from it accurate models for the distribution of clouds in redshift and H~I column density, constraining the opacity of the universe to H~I- and He~II-ionizing photons. We describe our assumptions about the physical properties of these absorbers and our models for the evolution and spectral shape of ionizing sources. In \\S~3, we describe the results of our radiative transfer calculations. These give the relation between the emitted spectrum of sources and the spectrum seen by the absorbing matter, as well as the implied opacity of the universe to He~II-ionizing photons. Using these models, we discuss the ionization conditions necessary to reproduce the mean levels of H~I and He~II Ly$\\alpha$ absorption seen in the current observations. In \\S~4, we discuss the fluctuations in the He~II absorption expected from variations in the number density of lines and the metagalactic radiation field. In \\S~5, we show how the effects discussed in this paper could be studied by future ultraviolet observations of He~II absorption. ", + "conclusions": "We have used a radiative transfer code and a new absorption-line sample to relate the spectrum of ionizing sources in the early universe to the ionization conditions in the \\Lya\\ forest and diffuse IGM. We find that a model with only line-blanketing and ionization by quasars is sufficient to explain the observed level of absorption, although contributions from stars or the presence of a diffuse IGM would help to explain the data. We show that the H~I and He~II absorption really does probe the metagalactic background radiation, and that the observed quasars are probably sufficient to ionize both H~I and He~II by redshift $z \\sim 3.5$. By considering fluctuations both in the density of clouds and the ionizing background, we find that a long redshift path is required to obtain the mean level of absorption. This is not entirely a drawback, however, as these fluctuations carry interesting information in themselves. Observations by FUSE, STIS, and future high-throughput spectrographs should allow us to determine the spectral shape and sources of the ionizing background with high accuracy and to disentangle the properties of the ionizing sources from those of the absorbers, leading to a much better understanding of the latter." + }, + "9802/astro-ph9802070_arXiv.txt": { + "abstract": "Ground and in flight calibrations of the MECS experiment on board Beppo-SAX have demonstrated that this is currently the best X-ray imaging experiment above 3 keV. The MECS on-axis PSF has a half power radius of about 1 arcmin. Moreover due to a fortunate combination of detector and mirror PSFs the total PSF depends only weakly on the energy. Finally the degradation of the PSF with off axis angle is negligible within an off-axis angle of 10 arcminutes. Encouraged by these results we developed techniques to analyze galaxy clusters observed with Beppo-SAX. In this proceeding we quantify spectral distortions introduced by the energy dependent PSF when performing spatially resolved spectroscopy of the core of the Virgo cluster. ", + "introduction": "The Medium Energy Concentrator Spectrometer (MECS; Boella et al. 1997/a) is one of the four narrow field instruments on board the Beppo-SAX satellite (Boella et al. 1997/b). The MECS operates in the energy band 1.3-10 keV with a field of view of $28'$ radius. The MECS consists of three units each composed of a grazing incidence Mirror Unit (MU), and of a position sensitive Gas Scintillation Proportional Counter (GSPC).\\\\ The Point Spread Function of the MECS (PSF$_{\\rm MECS}$) is the convolution between the MU PSF and the detector PSF. The MU and detector PSFs are described, respectively, by a lorentzian function L(r) and a gaussian function G(r). Both the lorentzian and the gaussian functions are energy-dependent. Typically the detector PSF improves with increasing energy, whereas the MU PSF improves with decreasing energy. The detector PSF dominates the core of the PSF$_{\\rm MECS}$ ($r \\simlt 2'$) whereas the MU PSF dominates the wings of the PSF$_{\\rm MECS}$ ($ \\simgt 2'$).\\\\ The analytical expression for the on-axis PSF as given in Boella et al. 1997/a is: $${\\rm PSF_{MECS}}(r,E)=\\frac{1}{2\\pi\\Big[R(E)\\sigma^{2}(E)+\\frac{r^{2}_{l}(E)} {2(m(E)-1)}\\Big]}\\cdot $$ \\begin{equation} \\Bigg\\{R(E)\\exp\\Bigg(-\\frac{r^2}{2\\sigma^2(E)}\\Bigg)+\\Bigg[1+\\Bigg(\\frac{r} {r_l(E)}\\Bigg)^2\\Bigg]^{-m(E)}\\Bigg\\}, \\end{equation} where $R(E)$, $\\sigma(E)$, $r_{l}(E)$ and $m(E)$ are algebraic functions of the energy E.\\\\ The integral of the PSF over the entire plane is normalized to unity:\\\\ $2\\pi\\int_{0}^{\\infty}{\\rm PSF_{\\rm MECS}}(r,E)r\\,dr\\equiv1$.\\\\ We used eq. 1 to evaluate the $50\\%$ and $80\\%$ power radii ($r_{50}(E)$ and $r_{80}(E)$) as a function of E:\\\\ \\vskip -0.5truecm $$2\\pi\\int_{0}^{r_{50}(E)}{\\rm PSF_{\\rm MECS}}(r,E)r\\,dr=0.5,$$ \\begin{equation} 2\\pi\\int_{0}^{r_{80}(E)}{\\rm PSF_{\\rm MECS}}(r,E)r\\,dr=0.8\\,. \\end{equation} \\begin{figure}[htb] \\vspace{9pt} \\epsfig{figure=proc_1.ps, height=7.5cm, width=7.5cm, angle=-90} \\caption{ {\\rm PSF}$_{\\rm MECS}$ Power Radii vs. Energy. The solid line represents the 50 \\% power radius, r$_{50}$(E), the dashed line represents the 80 \\% power radius, r$_{80}$(E) (see eq. 2).} \\end{figure} As shown in fig. 1 $r_{50}(E)$ is always $<2'$ and decreases with increasing energy. This is because at radii $<2'$ the PSF$_{\\rm MECS}$ is dominated by the gaussian PSF of the detector, that improves with increasing energy. The power radius $r_{80}(E)$ does not vary strongly with energy (see fig. 1) because of the combined effect of the improvement of the detector PSF and degradation of the MU PSF with increasing energy. \\begin{figure}[htb] \\vspace{9pt} \\epsfig{figure=proc_2.ps, height=7.5cm, width=7.5cm, angle=-90} \\caption{ Ratios between convolved profiles vs. radius r. Solid line: Ratio R$_{1}$(r) between the convolved profile $\\tilde{I}$(r,3 keV) and the convolved profile $\\tilde{I}$(r,6 keV). Dashed line: Ratio R$_{2}$(r) between the convolved profile $\\tilde{I}$(r,9 keV) and the convolved profile $\\tilde{I}$(r,6 keV)} \\end{figure} \\begin{figure}[htb] \\vspace{9pt} \\epsfig{figure=proc_3.ps, height=7.5cm, width=7.5cm, angle=-90} \\caption{ Correction vectors for the Virgo cluster observed with the MECS. Dot-dot-dot-dashed line: correction vector V$_{1}$(E$_{i}$) for the region $0'-2'$. Dashed line: correction vector $V_{2}$(E$_{i}$) for the region $2'-4'$. Dash-dotted line: correction vector $V_{3}$(E$_{i}$) for the region $4'-6'$. Dotted line: correction vector $V_{4}$(E$_{i}$) for the region $6'-8'$. Solid line: correction vector $V_{5}$(E$_{i}$) for the region $8'-10'$. } \\end{figure} ", + "conclusions": "" + }, + "9802/astro-ph9802185_arXiv.txt": { + "abstract": "To predict the X-ray observables associated with the diffuse baryons in clusters of galaxies, we develop a new physical model for such a hot intra-cluster plasma. Our framework is provided by the hierarchical clustering cosmogony for the dark matter, and by the standard FRW or Lema\\^itre cosmologies constrained by cosmic ages. As to the plasma dynamics and thermodynamics we propose a semi-analytical approach based on {\\it punctuated equilibria}. This comprises the following blocks that we compute in detail: Monte Carlo ``merging histories'' to describe the dynamics of dark matter condensations on scales of order $1-10$ Mpc, and the associated evolution of the gravitational potential wells; the central {\\sl hydrostatic} disposition for the ICP, reset to a new equilibrium after each merging episode; conditions of shock, or of closely adiabatic compression at the {\\sl boundary} with the external gas, preheated by stellar energy feedbacks. Shocks of substantial strength are shown to prevail at the outskirts of rich clusters in a universe with decelerated expansion. From our model we predict the $L-T$ relation, consistent with the data as for shape and scatter. This we combine with the mass distribution provided by the canonical hierarchical clustering; the initial perturbation spectra are dominated by Cold Dark Matter but include enough baryons to account for the high abundance sampled by the X--ray clusters, and are COBE--normalized. Thus we predict the $z$-resolved luminosity functions, with the associated source counts and redshift distributions. We predict also the complementary contribution by the unresolved groups and clusters to the soft X-ray background. These results are compared with two recent surveys from ROSAT; one defines the local luminosity function over nearly three decades of $L$, and the other shows little or no evolution out to $z\\sim 0.8$. Our results confirm that the critical cosmology coupled with Standard CDM is ruled out by its overproduction of local clusters. On account of underproduction, instead, we rule out open cosmologies (the cheapest way to solve the baryonic crisis and to freeze evolution), except for a narrow range around $\\Omega_o =0.5$; even there, we find the consistency with the full data base to be hardly marginal. For the CDM cosmogony with $\\Omega_o=0.3$ but in flat geometry, we obtain acceptable fits. For the tilted CDM perturbation spectrum with high baryonic content in the critical universe, we obtain marginal consistency. The cosmogonical/cosmological sectors of the cluster history are independently testable by means of a lower bound to the evolved temperature distribution, as can be measured with SAX and XMM out to moderate $z$. Finally, we discuss the effective limitations of X-ray clusters and groups as cosmological signposts, and their brighter prospects toward the astrophysics of the ICP and the cosmogony of large, high--contrast structures. ", + "introduction": "Groups and clusters of galaxies constitute cosmic structures sufficiently close to equilibrium and with sufficient density contrast ($\\delta\\approx 2\\, 10^2$ inside the virial radius $R$) as to yield definite observables, and possibly to offer reliable signposts for cosmology. They stand out to substantial depths of space-time not only in the optical band, but even more in X--rays. This is because their gravitational potential wells, shaped by a dominant dark mass $M$, contain not only baryons condensed into stars but also a larger amount of {\\it diffuse baryons}. The latter, with densities $n\\sim 10^{-3}$ cm$^{-3}$ and virial temperatures $k\\,T\\sim GM\\,m_H/10 R\\sim 5$ keV in rich clusters, satisfy the plasma condition $kT/ e^2\\,n^{1/3}\\gg1$ exceedingly well; in fact, they do by a factor $10^{11}$, vastly larger than for the baryons inside the stars. Such a hot {\\it intra--cluster plasma} (ICP) emits powerful X-ray luminosities $L\\propto n^2\\,R_X^3\\,T^{1/2}\\sim 10^{44}$ erg/s by optically thin thermal bremsstrahlung from central regions of overall radius $R_X\\sim 1$ Mpc. The temperature directly probes the height of the potential well, with the baryons in the role of mere tracers; on the other hand, the luminosity with its strong dependence on density reliably probes the baryonic content. Statistically, a definite $L-T$ correlation is observed (albeit with considerable scatter), and this provides the crucial link to relate the X-ray luminosity functions with the statistics of the dark mass $M$ or with that of the corresponding $T$. But groups and clusters are intrinsically complex systems. To begin with, their dynamical history is marked by extensive, repeated {\\it merging} of clumps, both in the form of nearly isotropic accretion of small units, and in the form of a few large, anisotropic coalescence events. This is shown in detail by all N-body simulations (see, e.g., Schindler \\& Mueller 1993; Tormen 1997; Roettiger, Stone \\& Mushotzky 1997), and is increasingly indicated by the data (see, e.g., Zabludoff \\& Zaritsky 1995; Henriksen \\& Markevitch 1996; Jones et al. 1997). The timing of such a dynamical evolution (and specifically the present merging rate) is set by the {\\sl cosmological} framework. Adopting the homogeneous isotropic FRW or Lema\\^itre cosmologies, the expansion is parametrized by the Hubble constant $h$ (in units of 100 km/s Mpc), and by the density parameter $\\Omega_o$ (to which the baryons contribute the fraction $\\Omega_B$); a possible additional contribution is given by $\\Omega_{\\lambda}$ associated to the cosmological constant. Given the cosmological framework, the {\\sl cosmogony} (i.e., the process of structure formation) is treated in terms of the standard hierarchical clustering scenario (see Peebles 1993), where all structures form by gravitational instability of initial density perturbations $\\delta$ in the dark matter (hereafter DM). Important parameters are the shape of their power spectrum $\\langle |\\delta_k|^2\\rangle\\propto k^{n_p}\\,T^2(k)$ at the recombination, and the normalization measured at large scales by COBE/DMR (Gorsky et al. 1996). The transmission function $T(k)$ depends on the specific model assumed for the DM; given this, the amplitude $\\sigma_8$ (i.e., the normalization extrapolated down to the relevant scale $8\\,h^{-1}$ Mpc) and the effective spectral slope $n_e$ depend on $\\Omega_o$ and $h$. These parameters also affect the growth factor $D(z)$ of the perturbations, which enters the actual predictions for the mass distributions of the condensed clusters or groups (see Appendix A). Further complexity is added by the {\\sl physics} of the ICP, and this constitutes our main aim here. To now, the ratio of the ICP to DM, and specifically its central density $n_o$ and its effective radius are poorly understood. But the former is especially important, as the factor $n_o^2$ enters the luminosities and {\\it amplifies} the observed variance. Observational information on the ICP and on the underlying DM dynamics is provided by the local $L-T$ correlation. At higher $z$, further statistical information is provided by the evolution of the X-ray luminosity function $N(L,z)$, or by its integrals like the source counts or the $z$-distributions. The problem is that the predictions of observables involve not only the ICP physics and the cosmogony (with their intrinsic variances) but also cosmology again (with its uncertainties), and so the various aspects are not easily disentangled. The sharpest result obtained so far rules out the attractively simple assumption by Kaiser (1986) that the ICP amount be proportional to the DM's from groups to clusters at all $z$ and $M$. A large number of subsequent works (of which we cite here the recent Mathiesen \\& Evrard 1997, Kitayama \\& Suto 1997, and Borgani et al. 1998) dealt with the combinations of these three {\\it sectors}, namely, cosmology, cosmogony and ICP physics. But most of these papers, following the start by Cavaliere \\& Colafrancesco (1988), approached the problem by parameterizing the dependences of the ICP/DM ratio on $M$ and $z$, e.g., in the form $L\\propto M^{p}\\,(1+z)^{s}$. While the parameters $p$ and $s$ are constrained to some extent by the local $L-T$ correlation, nevertheless there still remains a substantial {\\it degeneracy} (Oukbir, Bartlett \\& Blanchard 1997) between ICP physics and cosmology/cosmogony. That is to say, different combinations of $h$, $\\Omega_o$, $\\Omega_{\\lambda}$, $\\Omega_B$, $\\sigma_8$, $n_p$, $s$, $p$ provide close fits to the observables. Conversely, when the ICP parameters are varied, different cosmologies appear to be preferred by the data; the trade off particularly concerns $\\Omega_o$ and $ s$, which directly govern the global and the ICP evolution, respectively. To go beyond such degeneracies a {\\it physical} model is needed for the ICP, including the above complexities. Such a model must include: the histories of DM halos with their hierarchical merging events; the infall of the gas with the ensuing compression and shocks; the disposition of the ICP in the potential wells; its conditions at the boundary with the surrounding environment, which is modulated by the large scale structures and by stellar preheating. A model accounting for all the above is missing so far. We stress that the simulations using advanced hydrodynamics coupled with N--body codes hardly reach at present enough dynamic range (as discussed by Bryan \\& Norman 1997) to describe DM and ICP over the full range from $\\sim 50$ Mpc associated with the large scale structures (which guide the ongoing mergers of DM halos), to the inner 50 kpc where the ICP yields a considerable contribution to $L$. On the other hand, such high-resolution simulations do not include yet the stellar preheating with its crucial effects on the LT relation. This motivates us to develop here a semi-analytical model which includes, though in a simplified form, the features listed above. We describe the cluster evolution as a sequence of {\\it punctuated equilibria} (PE); that is to say, a sequence of hierarchical merging episodes of the DM halos (computed with Monte Carlo simulations), associated in the ICP to shocks of various strengths (depending on the mass ratio of the merging clumps), which provide the boundary conditions for the ICP to re-adjust to a new hydrostatic equilibrium. We show that our PE model predicts density and temperature profiles and the $L-T$ relation for clusters and groups consistent with the recent data. We then use our PE to predict the counts of resolved sources $N(>F)$ for faint fluxes down to $F > 2\\,10^{-14}$ erg/s cm$^2$, now accurately measured by Rosati et al. (1997). We predict also the {\\it complementary} observable constituted by the contribution of the unresolved groups and clusters to the soft X-ray background (XRB). We predict the $z$-distributions, and finally the full $z$-resolved luminosity functions to be compared with recent and forthcoming surveys. The paper is organized as follows. In \\S 2 we present and discuss our approach to the ICP astrophysics. In \\S 3 we give the X-ray observables in the form suited to our hierarchical clustering computations; details are supplied in the Appendixes A, B and C. In \\S 4 we present the results from our approach. The final \\S 5 is devoted to discussions and conclusions. \\section {The punctuated equilibria for the ICP} The X-ray luminosity of a cluster with radial temperature profile $T(r)$ and density profile $n(r)$ is given by \\begin{equation} L\\propto \\int d^3 r\\,n^2(r)\\, T^{1/2}(r)~; \\end{equation} the integration is over the emitting volume $r^3\\leq R^3$. Expressing $R$ from $T\\propto M/R\\propto \\rho\\,R^2$, eq. (1) recast into the form \\begin{equation} L \\propto \\overline{[n(r)/n_1]^2}\\,T^2\\,\\rho^{1/2}~, \\end{equation} where $\\rho$ is the internal DM density, $n_1$ is the density just exterior to the cluster boundary, and the bar indicates the integration over the cluster volume normalized to $R^3$. Note that eq. (2) applies in the isothermal case; the corresponding expression for a polytropic ICP is given in Appendix B. The simplest approach to the ICP state is that adopted in the self--similar model (Kaiser 1986) where $n \\propto\\rho$ is assumed, independently of $T$; then $L\\propto T^2$ obtains from eq. (2). The result conflicts with the observed correlation for rich clusters which is close to $L\\propto T^3$ (Edge \\& Stewart 1991; Mushotzky 1994; Tsuru et al. 1996; Mushotzky \\& Scharf 1997). Also, when combined with the standard hierarchical cosmogony, the assumption yields unacceptable fits to the local luminosity function (see, e.g., Kitayama \\& Suto 1997). Finally, it would predict the clusters at higher $z$ to comprise not only denser DM, but also equally denser ICP; since bremsstrahlung depends on $n^2$, this would imply a strong positive evolution of $N(L,z)$, which is certainly {\\it not} observed. Rather, the analysis of the ROSAT Brightest Cluster Sample (Ebeling et al. 1997) and of the ROSAT Deep Cluster Survey by Rosati et al. (1997), extending and strengthening the data by Collins et al. (1997) and by Nichol et al. (1997), indicate no evolution out to $z\\approx 0.8$ for $L<3\\,10^{44}$ erg/s; earlier surveys (EMSS, Henry et al. 1992) suggested even a (marginal) negative evolution of the bright clusters. So to derive the true {\\it scaling} of $L$ with $T$ we need a closer analysis of the ICP disposition relative to the DM. We propose a new approach based upon two cornerstones: the profiles $n(r)$ and $T(r)$ are given by the {\\it hydrostatic} equilibrium with the gradually changing gravitational potential; their normalizations, and so the central density, are set by the conditions at the {\\it boundary} with the external medium. The equilibrium profile may be effectively represented with a polytropic relation starting from the cluster boundary at $r_2$, say: \\begin{equation} n(r)/n_1=g(T)\\,\\Big[{T(r)\\over T_2}\\Big]^{1/(\\gamma-1)}~. \\end{equation} Here $T_2$ is the temperature just interior to the boundary; we conveniently use $g(T)\\equiv n_2/n_1$ for the ratio of the interior to the exterior density, to include any shock discontinuity at the boundary. The appropriate values for $\\gamma$ will be discussed in \\S 2.1 and \\S 2.4. Actually, the ICP is reset to a {\\it new} equilibrium after each episode of accretion or merging of further mass. In our PE approach, the history of such episodes is followed in the framework of the hierarchical clustering by Monte Carlo simulations, as explained below. Two relevant {\\it limiting} forms of eq. (3) are constituted by by the ``shock'' model of CMT97, and by the ``adiabatic'' models of Kaiser (1991) and of Evrard \\& Henry (1991). In both the outer gas is expected to be {\\sl preheated} at $T_1\\lesssim 1 $ keV (Ciotti et al. 1991; David et al. 1993, 1995; Renzini 1997) by $z\\lesssim 2$, due to feedback energy inputs following star formation and evolution. Preheating temperatures $T_1\\gtrsim 0.1$ keV also would prevent the cooling catastrophe from occurring, see White \\& Rees (1978); Blanchard, Valls Gabaud \\& Mamon (1992). In point of fact, Henriksen \\& White (1996) find from X-rays evidence for diffuse cool gas at $0.5 - 1$ keV in the outer regions of a number of clusters. In the present context, preheating will inhibit the attainment of the universal baryonic density in gravitational wells with virial temperatures comparable to $T_1$. These limiting models differ in their treatment of the boundary conditions and of the merging histories. \\subsection{Shocks and hydrostatic equilibrium} The key boundary condition is provided by the dynamic stress balance $P_2=P_1+m_H\\,n_1\\,v_1^2$, relating the exterior and interior pressures $P_2$ and $P_1$ to the inflow velocity $v_1$ driven by the gravitational potential at the boundary. We expect the inflowing gas to become supersonic in the vicinity of $R$, when $m_H\\,v_1^2> 2 kT_1$. In fact, many hydrodynamical simulations of loose gas accretion into a cluster (from Perrenod 1980 to Takizawa \\& Mineshige 1997) show shocks to form, to convert most of the bulk energy into thermal energy, and to expand slowly remaining close to the virial radius for some dynamical times. So we take $r_2\\approx R$ (which follows the structure growth, since $R\\propto M^{1/3}$), and focus on nearly static conditions inside, with $v_2^2<< v_1^2 $. The post-shock state is set by conservations across the shock not only of the stresses, but also of mass and energy, as described by the Rankine-Hugoniot conditions (see Appendix B). These provide at the boundary the temperature jump $T_2/T_1$, and the corresponding density jump $g$ which reads \\begin{equation} g\\Big({T_2\\over T_1}\\Big) = 2\\,\\Big(1-{T_1\\over T_2}\\Big)+\\Big[4\\, \\Big(1-{T_1\\over T_2}\\Big)^2 + {T_1\\over T_2}\\Big]^{1/2} \\end{equation} for a plasma with three degrees of freedom. Eq. (4) includes both {\\it weak} and {\\it strong} shocks. For weak shocks with $T_2 \\approx T_1$ (appropriate for small groups accreting preheated gas, or for rich clusters accreting comparable clumps), this converges to the truly adiabatic relationship $n_2/n_1= (T_2/T_1)^{3/2}$ up to second order inclusive, see Landau \\& Lifshitz (1959). On the other hand, it is shown in Appendix B that for strong shocks (appropriate to \"cold inflow\" as in rich clusters accreting small clumps and diffuse gas) the approximation $k\\,T_2\\approx -\\phi_2/3+3k\\,T_1/2$ holds, where $\\phi_2$ is the gravitational potential energy at $r_2\\simeq R$. Inside $R$, the temperature and density profiles $T(r)$ and $n(r)$ are matched to $T_2$ and to $n_2$ by {\\it polytropic} profiles or by their {\\it isothermal} limit. We numerically compute such profiles using the hydrostatic support of pressure against gravity (see Appendix B); for definiteness, we use the Navarro et al. (1996) representation for the potential and for the velocity dispersion (which varies slowly in the relevant region). Let us consider first for reference the simple analytical approximation provided by the standard isothermal model \\begin{equation} n(r)/n_2=[\\rho (r)/\\rho_2]^\\beta \\, , \\end{equation} (Cavaliere \\& Fusco-Femiano 1976) with the canonical exponent $\\beta \\equiv \\mu m_H \\sigma/kT_2$; here $\\sigma$ is the one-dimensional velocity dispersion at $R$, and $\\mu$ is the average molecular weight. For the purpose of the analytical approximation we may use the strong shock limit for $T_2$, and find numerical values for $\\beta$ ranging from about $ 0.5$ for groups to 0.9 for rich clusters (a trend consistent with data being collected by M. Girardi and coauthors, private communication). This implies that $R_X/R$ is larger in the former than in the latter (see CMT97). The full numerical computations using the expression for $T_2$ given in Appendix B, and $\\sigma (r)$ and $\\phi (r)$ from Navarro et al. (1996) confirm this trend and give results shown in fig. 1. Note that even in the isothermal case the emission-weighted temperature (integrated along the line of sight) declines outwards but only very slowly. The observed stronger decline requires a polytropic equilibrium, where the run of $T(r)$ steepens with the index $\\gamma$ increasing from 1. In the Appendix B we recall the basic relations, and show that the variations induced in the volume-averaged luminosity by increasing $\\gamma$ are small. Thereafter, we adopt the value $\\gamma=1.2$, which also yields for rich clusters an integrated baryonic fraction 0.15 out to $R$. The result for the emission-weighted $T(r)$ (see fig. 1) is a mild decrease out to $r\\sim$ 1 Mpc in agreement with the observations (Hughes, Gorenstein \\& Fabricant 1988; Honda et al. 1997; Markevitch et al. 1997), followed by a sharper drop as indicated by state-of-the-art simulations (e.g., Bryan \\& Norman 1997). Note that in fig. 1 the discontinuity at the shock has been smeared out to a smooth drop by the finite resolution, taken at 100 kpc for comparison with the simulations and with the forthcoming observations. \\subsection{Merging histories and the $L$-$T$ correlation} The luminosity $L\\propto g^2~ \\overline{n^2(r)/n_1^2}~ T^2 \\; \\rho^{1/2}$ is statistically affected by the {\\it merging histories} as follows. For a cluster or group of a given mass (or temperature), the effective compression factor squared $\\langle g^2 \\rangle$ is obtained upon averaging eq. (4) over the sequence of the DM merging events; in such events, $T_2$ is the virial temperature of the receiving structure, and $T_1$ is the higher between the stellar preheating temperature and that from ``gravitational'' heating, i.e., the virial value prevailing in the clumps being accreted. All that is accounted for in our model using Monte Carlo simulations of the hierarchical growth of the DM halos; these are based on merging trees corresponding to the excursion set approach of Bond et al. (1991), consistent with the Press \\& Schechter (1974) statistics (see CMT97). The averaging procedure is dominated by the events occurring within the last few dynamical times; it results in lowering $\\langle g^2 \\rangle$ compared to $g^2$, because in many events the accreted gas is at a temperature higher than the minimum preheating value $T_{1\\ell}\\approx 0.5 $ keV. In addition, an intrinsic {\\it variance} is generated, reflecting and amplifying the variance intrinsic to the merging histories. The net result is shown in fig 2, and commented in its caption. In agreement with the observations (Edge \\& Stewart 1991; Mushotzky 1994; Tsuru et al. 1996; Ponman et al. 1996), the {\\it shape} of the average $L-T$ relation flattens from $L\\propto T^5$ at the group scale (where the nuclear energy from stellar preheating competes with the gravitational energy from infall) to $L\\propto T^3$ at the rich cluster scales. At higher temperatures the shape asymptotes to $L\\propto T^2$, the self-similar scaling of pure gravity. Notice the intrinsic {\\it scatter} due to the variance in the dynamical merging histories, but amplified by the $n^2$ dependence of $L$. The average normalization formally rises like $\\rho^{1/2}(z)$, where $\\rho$ is the effective external mass density which increases as $(1+z)^2$ (Cavaliere \\& Menci 1997) in filamentary large scale structures hosting most groups and clusters (see Ramella, Geller \\& Huchra 1992). This implies a factor 1.3 at $z=0.3$, consistent with the observations by Mushotzky \\& Scharf (1997). Further weakening of the $z$-dependence will comes from the increasing depth of the central $\\phi_o$ for distant structures of given $M$, as predicted by Navarro, Frenk \\& White (1996). \\subsection {The adiabatic models} At the other extreme, the models by Kaiser (1991), and Evrard \\& Henry (1991) obtain from PE under two limits, appropriate only for rapidly expanding universes, as we discuss below. The first limit correspond to no currently active merging, with shocks moving outward and vanishing. In such conditions, at the boundary $T_2 \\approx T_1$ holds; with $T_1$ staying nearly constant after the dynamical freeze out, this implies $g \\equiv n_2 /n_1 \\approx 1$. Correspondingly, the central density scales approximately as $n_o\\propto n_1\\,(T_o/T_1)^{1/(\\gamma-1)}$. The ``adiabatic'' models require also a second limit, concerning the internal gas distribution. The value $\\gamma = 5/3$ is taken at the center, but an isothermal $\\beta$ profile is assumed (with a fixed $\\beta$), based on a King--like DM distribution. A constant baryonic fraction at $R$ is then required, and this forces the core radius to scale as $r_c \\propto M^{1/3-1/3\\beta}$. Thus $L\\propto T^{2+3(2-1/\\beta)/2}$ obtains, with the normalization $\\rho^{1/\\beta -3/2}(z)$. Finally, the value of $\\beta$ is chosen as an input. The choice $\\beta =2/3$ for both clusters and groups leads to the model of Evrard \\& Henry (1991), in which $L \\propto T^{11/4}$ obtains, with constant normalization. The choice $\\beta =1$ leads to the somewhat different model of Kaiser (1991), in which $L\\propto T^{3.5}$ obtains, with the normalization anti-evolving like $(1+z)^{-3/2}$; this will become even more negative when the evolution of $\\phi$ is taken into account, so conflicting with the data of Mushotzky \\& Scharf (1997). \\subsection{ICP models and cosmology} Shock and adiabatic models can be characterized in terms of entropy (see Bower 1997). Actually, the modes of entropy production and distribution correlate with the global dynamics. Collapses, merging and the induced shocks are currently ongoing in the {\\it critical} universe, so that strong shocks form close to the virial radius. Entropy is continuously generated in the outer regions, so that its radial distribution is raised outwards. Then the effective $\\gamma$ will be close to one, leading to a roughly flat $T(r)$ inside the shock. The density is determined by the boundary conditions after eqs. (4). Shocks are weaker in groups, the density profile is shallower, and the $L-T$ relation steeper. Conversely, in an {\\it open} universe most dynamical action is moved back to early times: merging and mixing occurred early on, and then subsided; shocks had time to expand beyond $R$ and weaken; correspondingly, the accretion petered out under nearly adiabatic conditions for groups and for clusters as well. Then the effective $\\gamma$ is closer to $5/3$, and this may be used to roughly scale the central densities with different virial temperatures $T$ to obtain $L\\propto T^{3.5}\\,R_X^3$. The two adiabatic models adopt additional, and different, assumptions concerning $r_c$ or $R_X$, i.e., $r_c\\propto T^{-0.25}$, or $R_X=const$, as discussed in \\S 2.3. For open cosmologies with $\\Omega_0\\approx 0.5$, or for flat ones with $\\Omega_{\\lambda}=1-\\Omega_0$, the present deceleration is {\\sl intermediate} between the two above cases, and the applicability of the shock or of the adiabatic model is not so clearcut; we shall consider both, finding similar results as is expected. ", + "conclusions": "In this paper we have computed the X-ray observables for groups and clusters of galaxies. As anticipated in the Introduction, we use -- rather than continuous and possibly degenerate parametrizations -- only discrete combinations of {\\it physical} models appropriate for the Dark Matter and for the Intra-Cluster Plasma. We first list our results, and then discuss them in detail. We have developed the punctuated equilibria (PE) model for the ICP state and dynamics. This is comprised of the following two components. As for {\\it single} clusters, we have used a polytropic $\\beta$-model which yields temperature profile $T(r)$ (see fig. 1) in good agreement with the observations. We predict the ICP density profile $n(r)$ and the brightness profile to be flatter for groups than for clusters, corresponding to a larger extension of the ICP relatively to their gravitational radii. As for {\\it statistics}, we convolved the ICP equilibria with the histories of DM halos, and predicted the $L-T$ correlation to take the form shown in fig. 2, in agreement with the data. In addition, we predicted an {\\it intrinsic} variance with the minimum value also represented in fig. 2. Based on our PE model we then proceeded to compute for various standard cosmological frameworks the local and the evolved {\\it luminosity functions} of galaxy clusters, that we compared with the data (fig. 3a, 4, 5a, 6a). We derived also the number counts (fig. 3c, 5c, 6c), the $z$-distributions (fig. 9) and the contribution to the soft X-ray background (fig. 3d, 5d, 6d). Our results are summarized in figs. 7a and 7b; these show that the set of acceptable cosmogonies/cosmologies is restricted to three disjoint {\\it domains}: $0.4<\\Omega_o<0.5$ for the standard CDM; $\\Omega=1$ for the Tilted CDM; $\\Omega_o\\approx 0.3$ for CDM in flat geometry. In fig. 9 we summarize the {\\it confidence} levels at which the data are matched. We next proceed to discuss in detail the results listed above. \\subsection{ICP state in evolving DM halos} The ICP {\\it state} in the hierarchically evolving gravitational wells constitutes the focus of our new approach. We propose that such state follows suit, passing through a sequence of {\\it punctuated equilibria} (PE) that we compute semi-analytically. These computations comprise: the merging histories of the DM potential wells, obtained with a large statistics from Monte Carlo simulations of the hierarchical clustering; the inner {\\sl hydrostatic} equilibrium disposition, updated after each merging episode; and the {\\sl boundary} conditions provided by strong and weak shocks, or even by a closely adiabatic compression, depending on the ratio of the infall to the thermal energy in the preheated external medium. The results of our model depend on two parameters, the external temperature $T_1$ and density $n_1$, which are {\\sl not} free. Specifically, we use for $T_1$ the lower bound $T_{1\\ell}=0.5$ keV provided by the literature on stellar preheating; in the merging events the effective $T_1$ is the virial temperature of the incoming clumps, when this is larger than 0.5 keV. The value of $n_1$ for rich clusters is related to the DM density by the universal baryonic fraction $\\Omega_B\\approx 0.15$. Note that our PE model does not require strict spherical symmetry, but rather that the residual internal velocities be smaller than the inflow velocity. So they can include merging episodes ranging from nearly isotropic accretion of small clumps and diffuse gas, to anisotropic coalescence of comparable clumps along filaments of the large scale structures. The expression of the bolometric luminosity is proportional to $g^2=(n_2/n_1)^2$, the square of the density jump at the bounding shock. The average of such factor over the merging histories is what gives to the statistical $L-T$ correlation the {\\sl curved} shape shown in fig. 2. For rich clusters we obtain $L\\propto T^3$. This flattens to $L\\propto T^2$ for larger $T$, corresponding to the saturation of the shock compression factor, i.e., $g(T/T_1)\\rightarrow 4$ when $T\\gg T_1$. At the other end, the correlation steepens to $L\\propto T^5$ in the group range, where $T/T_1\\sim 1$ and the shocks are substantially weakened by the preheating temperature in the infalling clumps. The amplitude of the $L-T$ correlation rises gently proportionally to $(1+z)\\propto \\rho_1(z)^{1/2}$ where $\\rho_1$ is the density in the large scale structures hosting clusters and groups. In addition, our PE approach predicts an intrinsic {\\it variance} of dynamical origin due to the different merging histories, and built in the factor $g^2$. Such variance constitutes a lower bound, in view of additional contributions from the variance in the ambient density, and from the central luminosity associated with cooling flows (Fabian et al. 1994; White, Jones \\& Forman 1997). \\subsection{Contact with hydrodynamical simulations and with observations} The PE model includes, in a simplified semi-analytical form, compression and shocks at the boundary with the surrounding environment, which is modulated in density by the large scale structure and in temperature by the stellar preheating. Simulations now clearly show the {\\sl shocks} occurring also in major merging events (Schindler \\& Mueller 1993; Roettiger, Stone \\& Mushotzky 1997). The inclusion of the Rankine-Hugoniot conditions rises the internal temperature at the expenses of the inflow velocities. Complex features are found like residual kinetic pressures, and unmixed hot spots in the temperature distribution; but over times of about 2 Gyr the residual kinetic pressure over cluster scales reduces to less than 20 \\% of the thermal one. On the other hand, in the hierarchical clustering such major events are rare; our Monte Carlo simulations give a probability $\\lesssim 20 \\%$ for very asymmetric events with mass ratios 2:1 or larger occurring within 2 Gyr from the cluster observation. In addition, in such events the ICP temperature in the infalling subcluster is comparable to the virial value in the main cluster. Such major events with their low frequency and large $T_1$ yield a minor contribution to the {\\sl statistical} $\\langle (n_2/n_1)^2\\rangle$. Our semi-analytical model describes only crudely these transient if conspicuous features, to focus on the lesser and more symmetric events which contribute the most to the $L-T$ relation. At the extreme of spherical accretion of loose gas the simulations (see Takizawa \\& Mineshige 1997) show in detail that {\\sl shocks} also form and expand slowly, to leave inside a declining temperature profile and a steeper density profile. Our PE model yields temperature profiles {\\sl decreasing} as shown in fig. 1. These agree with the published data (see Hughes et al. 1988; Honda et al. 1997; Markevitch et al. 1997); they agree also with the results from state-of-the-art simulations (Bryan \\& Norman 1997) obtained by running on supercomputers advanced 3D Eulerian codes with adapting mesh and reliable shock capturing methods. While the high-resolution simulations are limited (for now and for some time to come) to the condition of no stellar preheating suitable only for very rich clusters, our model includes the effects of stellar preheating of the outer gas over the whole range from groups to clusters. \\subsection{Constraining cosmology} With the ICP state so described, we proceeded to constrain the cosmological parameters. After the observations by Rosati et al. (1997), the main rule of the game turns out to be as follows: the dynamical evolution contained in the standard Press \\& Schechter formula (eq. 9) must {\\sl combine} with the evolution of the $L-T$ correlation and with cosmology to yield closely non-evolutionary $N(L,z)$. We stress that such combinations are severely selected in our approach. In fact, strong {\\sl shocks} are common in the {\\it critical} cosmology, where accretion and merging activity are currently ongoing, and there {\\sl shocks} apply in full; these include also weak shocks for small groups with virial temperatures below 1 keV, and for those rich clusters which merge with comparable clumps. Closely {\\sl adiabatic} compressions, instead, prevail for all structures in very {\\it open} universes with high formation redshifts, and little or no strong shocks and mixing at present. This defines the domain of applicability of the two adiabatic models by Kaiser (1991) and Evrard \\& Henry (1991). In detail, our results are as follows (see fig. 7). We confirm that Standard CDM in the {\\it critical} cosmology is definitely ruled out on account of its overproduction of local clusters and of their considerable positive evolution. We also rule out {\\it open} cosmologies with $\\Omega_o\\leq 0.3$, the simplest way to enforce little evolution and also to provide a solution to the baryonic crisis. In fact, these cosmologies when COBE-normalized yield a severe deficit in the local luminosity function and in the source counts. Thus, we investigated more {\\it elaborate} solutions: the critical universe with tilted CDM and a high baryon content; and two cases of intermediate deceleration, comprising CDM with canonical nucleosynthesis either in mildly open universes with $\\Omega_o\\approx 0.5$, or in flat geometry with $\\Omega_o=0.3$ and $\\Omega_{\\lambda}=0.7$. With the first solution we obtain (marginally) acceptable results (see figs. 3a-3c) for the local luminosity function $N( L,0)$, for the source counts, and for the contribution to the soft XRB, within the present uncertainties of the data and within the variance intrinsic to the theory. We note that the cosmological/cosmogonical sectors by themselves may be tested independently, based on the fast evolution with $z$ (see fig. 3b) of the temperature distribution in the critical case, as pointed out by many authors (see Oukbir \\& Blanchard 1992; Henry 1997); this constitutes an important program for the satellites SAX and XMM As for comologies with {\\it intermediate} deceleration, here neither the PE nor the adiabatic models for the ICP are cogently indicated; thus we considered both, obtaining generally similar results as expected. For open cosmologies in particular, the results are inconsistent with the observations of the local luminosity function and of the counts, except for the range $\\Omega_o=0.45 \\div 0.55$; even there the counts are excessive at more than the formal $99$\\% confidence level (see fig. 9), the excess being larger for the adiabatic models. The excess is due to built--in reasons, that is, the relative large amplitude $\\sigma_8$ and the relatively steep shape of the counts, as spelled out in \\S 4.4. Instead, in the $\\Omega_o +\\Omega_\\lambda= 1$ cosmology a manageable count excess is obtained. We note that the Kaiser's (1991) variant of the adiabatic models does not yield such an excess for $\\Omega_o\\approx 0.5$, due to its normalization decreasing at high $z$. However, the local luminosity function computed from this model overestimates the number of brightest clusters, due to the strong dependence $L\\propto T^{3.5}$. Moreover, the normalization decrease at high $z$ is hardly consistent with the data by Mushotzky and Scharf (1997). Finally, it yields a fast, negative evolution of $N(L,z)$ barely consistent with the data in the survey by Rosati et al. (1997) (the deficit is truly fatal in the critical or in the flat cosmology). As luminosities larger than some $3\\, 10^{44}$ erg sec$^{-1}$ are little represented in that survey, a strong test for such a negative evolution concerns any deficit at bright fluxes in the redshift distribution from a large--area survey. So we show in fig. 10 the redshift distribution computed also for this model. We have conservatively chosen to focus on a limiting flux $F = 4\\, 10^{-14}$ erg sec$^{-1}$ for which the sky coverage is nearly $100$\\% (P. Rosati, private communication), and any incompleteness is out of question. Incompleteness due to surface brightness may be relevant at fainter fluxes, depending on the cluster and group profiles. We plan to treat such issue elsewhere, but here we point out that in our approach the impact of any such incompleteness is limited by the {\\sl complementarity} between counts of resolved sources and contribution to the XRB from the rest. Our summary is that many combinations of standard cosmogonies/cosmologies with ICP models are ruled out. A relatively {\\it small} set of disjoint cosmologies/cosmogonies survive, as shown by fig. 7: $\\Omega=1$ with tilted CDM and high baryonic abundance combined with the Punctuated Equilibria, which is marginally consistent with the data; CDM in open cosmology with $0.5< \\Omega_o<0.55$, which is barely consistent using the PE, and even less so using the adiabatic models; CDM with $\\Omega_o=0.3$ and $\\Omega_{\\lambda}=0.7$, which is consistent using either the PE and the adiabatic models. So cosmological parameters can be constrained on the basis of X--ray clusters, but only {\\it up to a point}; for example, the residual uncertainty in the density parameter is $\\Delta\\Omega_o/\\Omega_o>20 \\%$. \\subsection{What next} To what extent enlarging the data base on X--ray clusters will help in further constraining cosmology? Here we argue that the variance intrinsic to the hierarchical clustering, and amplified by the ICP emissivity, sets an effective {\\it limitation}. In fact, fig. 7 shows that the present Poissonian error bars in the observed faint counts are already smaller than the (minimum) intrinsic variance in the predicted ones. Decreasing the former with richer, faint surveys will hardly provide a sharper insight into cosmology unless one reduces both the uncertainty concerning $\\sigma_8$ and the larger one concerning $L_{o}$; in fact, the two enter with comparable weights eq. (10), since it is seen that $\\Delta \\sigma_8/ \\sigma_8$ acts roughly as $(n_{e}+3)\\Delta L_{o}/6 L_{o}\\approx 0.2 \\Delta L_{o}/L_{o}$. To what point are these reductions feasible? On the theoretical side, the minimum $\\Delta L_{o}/L_{o}$ of dynamical origin may be sharpened by Monte Carlo simulations so extensive as to provide the full scatter distribution. But then one must tackle also the enhanced emissivity produced or signaled by cooling flows, correlated with higher ambient densities; this we shall treat elsewhere (Cavaliere Menci \\& Tozzi, in preparation). On the observational side, one needs a large statistics for the distributions of $L$ and $T$; this will help in deriving narrower $L-T$ correlations for subsamples categorized in term of mass deposition rates from cooling flows, see White, Forman \\& Jones (1997). Such aim calls for spectroscopic measurements of $T$, which are obviously harder than the bolometric $L$, and require SAX or even XMM. However, we stress that such efforts will find soon a more {\\it proper} aim than constraining $\\Omega_o$. This is because soon MAP (Bennett et al. 1997), and subsequently PLANCK (Bersanelli et al. 1996), will accurately measure on very large and still linear scales not only the perturbation power spectrum (from which $\\sigma_8$ is derived), but also directly $\\Omega_o$ to better that 10\\%; this will supersede constraints set at cluster scales gone {\\it non-linear}. Once the cosmological framework has been fixed, the study of groups and clusters in X-rays will resume what we submit to be its proper course; that is, the physics of systems of intermediate {\\it complexity} which is comprised of the DM and of the ICP component. With the latter fully understood and the scatter in the $L-T$ relation assessed, cluster X--raying will finally expose the underlying process of {\\it non-linear} condensation of DM on scales $1-10$ Mpc. Then any mismatch concerning the number counts or $N(L,z)$ will be telling of failures either in the CDM spectra or in the current representation of cosmogony in terms of the Press \\& Schechter formula. As a relevant example, we recall from \\S 4.4 that even the acceptable models we computed tend to exceed the observed faint counts, and can be brought to consistency only at the lower end of the current uncertainty concerning $\\sigma_8$. On the other hand, the corrections to the Press \\& Schechter formula currently discussed yield a larger number of clusters. For example, Jain \\& Bertschinger (1995), and Gardner, Tozzi \\& Governato (1998) find that the threshold $\\delta_c$ must be lowered from the canonical value 1.69 to 1.5, at least at $z\\geq 1$ if not already at $z=0$; a similar trend obtains considering that the formation redshift is always greater than the observation's as discussed in \\S 4.4. If MAP will provide definite values of $n_p$, $\\sigma_8$ and $\\Omega_o$ such as to enhance the excess in the faint counts, then the Press \\& Schechter rendition of the non-linear cosmogony will have to be reconsidered." + }, + "9802/astro-ph9802117_arXiv.txt": { + "abstract": "We have identified a candidate $\\sim 1-2 \\times 10^5$ year old luminous white dwarf in NGC 1818, a young star cluster in the Large Magellanic Cloud. This discovery strongly constrains the boundary mass $M_c$ at which stars stop forming neutron stars and start forming white dwarfs, to $M_c \\gtsim 7.6 \\msun$. \\bigskip \\noindent {\\bf Key words:} Stars:white dwarfs -- globular clusters: individual: NGC 1818 ", + "introduction": "Stars of mass comparable to that of the sun evolve to form white dwarfs while, above some critical mass, $M_c$, stars detonate as type II supernovae instead, leaving neutron stars as remnants. The determination of $M_c$ is important for understanding stellar evolution. It also has implications for the chemical evolution of galaxies, in that it affects the rate of enrichment of the interstellar medium and the total population of pulsars. The relationship between the mass of a white dwarf and that of its progenitor, as a function of the metallicity of the progenitor, is also important for understanding stellar evolution in the early universe (\\refto{Jeff,vdB}). Predictions of $M_c$ range from $6-10\\, \\msun$ depending on the details of models and on the metallicity of the star (\\refto{Weid,Jeff,gar97,Hur}). Almost no observational constraints on $M_c$ are available except from one Galactic open cluster in which the most luminous of four white dwarfs implies $M_c > 5.7 \\, \\msun$ (\\refto{Koe}). Statistical extrapolations of these data suggest $M_c \\sim 8 \\, \\msun$ but with large formal uncertainty (\\refto{Jeff}). The young rich star clusters in the Large Magellanic Cloud (LMC) are particularly well suited for studies of the evolution of intermediate mass stars. They typically contain an order of magnitude more stars than Galactic open clusters, and membership determination is not generally problematic as it is with Galactic clusters, which are often superposed against a dense curtain of disk stars. The clusters with age $\\sim 2-4 \\times 10^7$ yr have main-sequence turnoff masses $\\sim 7.5 - 9.5 \\msun$, (\\refto{Will}) interestingly close to the theoretical range of values for $M_c$. Models show that turnoff stars of this mass spend $\\sim2\\times 10^6$ yr on the red giant branch (\\refto{Hur,Scha92}) before ejecting their hydrogen envelope and leaving rapidly fading white dwarfs (assuming they in fact have masses $< M_c$). Rich LMC clusters typically have $\\sim 10-20$ red giant branch stars, so we would expect the youngest remnant star to be $\\sim 1-2 \\times 10^5$ years old. The models of Wood give luminosities for white dwarfs of this mass and age of $30-100 L_\\odot$, with temperatures of $\\sim $100,000 K (\\refto{Woo92,Woo2,dan}). Cooling is dominated by neutrino emission, and the remnants fade proportionally to $\\sim t^{-2.5} $ or faster, depending on model details. Adopting a distance modulus for the LMC of $(V-M_V)=18.5$ (\\refto{Pan91}) would imply apparent magnitudes for these young white dwarfs of $V\\sim 18-19.5$. They would be visible even from the ground, although because of crowding, Hubble Space Telescope (HST) observations would be required for accurate photometry. The white dwarfs would be distinguished from comparably bright main-sequence stars by their extreme blueness. To investigate the possibility of finding young white dwarfs in LMC clusters, we analysed HST archive images of NGC~1818. This cluster has mass $\\sim 2.8 \\times 10^4 \\msun$, core radius $r_c=2.0$ pc, half-mass radius $r_h \\sim 14$ pc (\\refto{Els87}), and age $\\sim 2-4 \\times 10^7$ yr (\\refto{Will}). There are currently $\\sim 16$ stars in the red giant phase. \\section {Observations} The images of NGC 1818 were obtained with the Wide Field and Planetary Camera (WFPC2) on 1995 December 8, with the F336W ($\\equiv U_{336}$), F555W ($\\equiv V_{555}$), and F814W ($\\equiv I_{814}$) filters. Total exposure times are 960, 880 and 1290 seconds respectively (\\refto{Hun97}). The images in each filter were coadded with a median filter to eliminate cosmic rays. DAOFIND was used to detect objects $4 \\sigma$ above the background. Point-spread function (PSF) fitting was used to help eliminate spurious detections: these include primarily structure in the PSF, particularly around saturated stars, and bright pixels along diffraction spikes. The final photometry was performed using an aperture with radius 2 pixels, and aperture corrections and zero points were applied (\\refto{Hola,Holb,Els98}). A value of $E(B-V)=0.05$ was adopted for the reddening. A color-magnitude diagram (CMD) in $V_{555}$ vs $(U_{336}-V_{555})$ for the Planetary Camera (PC) revealed a prominent sequence of binary stars, and we explore the implications for the binary population of NGC 1818 elsewhere (\\refto{Els98}). CMDs for the three WFC chips combined are shown in Figs. 1a and b. These cover a radial range from $\\sim 2$ core radii, out to about half way to the edge of the cluster. Stars with $V_{555} < 17.5$ are saturated; the main-sequence turnoff is at $V_{555} \\approx 14$. (We note the presence of an apparent gap in the main-sequence at $V_{555}\\sim 20.3$ which may be due to a possible jump in stellar magnitudes around $\\sim 1.5 - 2 \\msun$ due to the onset of convective overshooting (\\refto{Hur}).) Models suggest that white dwarfs should have $(U-V)_0 \\sim -1.5$ and $(V-I)_0 \\sim -0.4$ (\\refto{Woo2,Che93}); these values are roughly independent of age and metallicity over the range of interest. For white dwarfs the model colors are equivalent to $(V_{555}-I_{814})=-0.44$ and $(U_{336}-V_{555}) = -1.85$ (\\refto{Holb} Fig. 11). The boxes drawn in Figs. 1a and b delimit plausible ranges of colors and magnitudes for a young white dwarf. We inspected visually all detections within these boxes, and in a similar range of magnitudes and colors for the CMDs for the PC chip. All but one of the objects turned out to be spurious detections as described above. One, however, turned out to be unambiguously stellar. It is indicated with a filled circle in the CMDs, and is circled on the image of the cluster shown in Figure 2. Could the photometry, or identification as a white dwarf, be in error? Poisson error bars are included in Fig. 1b, and indicate that the extreme blueness of this object cannot be due to random photometric errors. Nor is it likely that the colors are biased bluewards by a residual cosmic ray in the F336W image: six independent images were coadded, so a cosmic ray would have had to hit the same pixel in several of them. Goodness-of-fit parameters (DAOPHOT's `sharpness' parameter and $\\chi^2$) derived from fitting a PSF to this object are within the range expected for a star: there is no evidence that it is strongly peaked, as a cosmic ray would be, or resolved, as a background galaxy would be. (A background galaxy this bright would be easily identifiable in a WFPC2 image.) The object is extremely unlikely to be a quasar: quasars of this magnitude have surface density $\\sim 1$ per square degree (\\refto{Co,Boy}), so the probability of finding one in our field is $\\sim 10^{-3}$. It is even more unlikely to be a foreground star in the halo of our Galaxy: the number of Galactic stars of this magnitude expected in a WFPC2 field of view at high galactic latitude is less than one (\\refto{San}). At a distance of, say, 10 kpc, a Galactic white dwarf would have absolute magnitude $M_V \\sim -3.5$, which is far too bright to be an old white dwarf. On the other hand, given the number of stars on the giant branch, we would expect to find about one young white dwarf if the turnoff mass in NGC 1818 is indeed less than $M_c$. Nor should we be surprised to find it outside the core of the cluster: more than half the red giants in the cluster are outside the core, and four are further from the cluster center than our white dwarf candidate. If we see one white dwarf with $V \\sim 18.5$, would we expect to see many more fainter ones? This depends on the cluster mass function and on the difference between the turnoff mass and $M_c$. Given the number of red giants in the cluster, and assuming that $M_{turnoff} \\approx M_c$, we might expect there to be $2 \\pm 2$ other white dwarfs with magnitude up to $\\sim 4$ mag fainter than our candidate (ie. with $18.5 < V < 22.5$). Given that sample of stars in the PC image is incomplete due to the presence of many saturated stars, and that the WFC chips cover just over 50\\% of the outer parts of the cluster, the expected numbers are consistent with our observations. If $M_c$ is significantly greater than the turnoff mass of the cluster then stars with slightly smaller masses than the current turnoff mass will have evolved to form white dwarfs which by now will have cooled and faded. We might therefore expect there to be 4--8 other white dwarfs at fainter magnitudes still; the exact number expected is dominated by small number statistics when the mass function is extrapolated. Due to the rapid fading, we would expect these to be $\\sim 5$ magnitudes fainter (ie. $V \\sim 23.5$) and they would be difficult or impossible to detect in the currently available WFPC2 images. The luminosity function of white dwarfs expected in young LMC clusters is discussed further in \\refto{Els98}. Deeper images of NGC 1818 are scheduled during HST Cycle 7 (Project 7307) to look for a sequence of fainter, older white dwarfs. \\section {Discussion} What constraints on $M_c$ can we infer from the presence of a young white dwarf in NGC 1818? $M_c$ must be greater than the turnoff mass, so the main task is to determine an accurate value for this. The main uncertainties are not in determining the magnitude of the turnoff, but in converting this magnitude to a mass. This requires a knowledge of both the metallicity and of whether stellar evolution models with or without convective overshooting are more appropriate. The metallicity of NGC 1818 has been determined from spectra of two stars to be [Fe/H]$= -0.8$ (\\refto{Will}). This value is, however, lower than the expected value for the young population in the LMC (\\refto{Olsz}), which is thought to be more like $-0.2$. The membership of one of the stars is doubtful, and $-0.8$ may in fact be too low. Two models with convective overshooting, and metallicities [Fe/H]$= -1.3$ and $-0.4$ give an age for NGC 1818 of $4 \\times 10^7$ yr, and a turnoff mass 7.6$\\pm 0.1 \\msun$. Two models without overshooting give an age $2.2 \\pm 0.2 \\times 10^7$ yr, and a turnoff mass 9.0$\\pm 0.5 \\msun$ (\\refto{Will,Scha92}). From our own stellar evoultionary models, which include convective overshooting, we predict $M_c=7.0$ for [Fe/H]$= -0.8$ and $M_c=7.7 \\msun$ for [Fe/H]$= -0.3$. The convective overshoot parameter for these models is $\\Lambda_c = 0.28$ for $M \\sim 7.5 \\msun$, which is close to but slightly higher than the models of Schaller \\etal used by Will \\etal ~Other groups use different values for $\\Lambda_c$, and a full discussion of the effect of the choice of this parameter on our results is deferred to a future paper (\\refto{Els98}). Our observations imply $M_c \\gtsim 7.6 \\msun$, and perhaps $M_c \\gtsim 9.0 \\msun$. A better determination of the metallicity of NGC 1818 would help constrain the value of $M_c$ further, as would deeper images in which any older white dwarfs would be visible. Another uncertainty affecting observational constraints on $M_c$ is the possiblity of an age spread among the stars near the main-sequence turnoff. For example, it is possible that the lower mass stars formed first, by a sufficient margin that some stars slightly below top of the main-sequence are evolving on to the red giant branch at the same time as more massive stars. The most massive stars have main-sequence lifetimes about 5 Myr, so the total age spread is at most a few Myr. Such a spread in age would mean that stars evolving on to the red giant branch could have masses $\\ltsim 0.3 \\msun$ less than the measured turnoff mass which would imply $M_c \\gtsim 7.3 \\msun$. Measuring such an age spreads requires accurate photometry for stars just below the turnoff, and we will be acquiring this during Cycle 7. Another source of uncertainty is the possibility that the object is a member of a mass--transfer binary, but this is unlikely and would require it now be a tight white dwarf -- neutron star binary. Because young white dwarfs are bright and very blue, detecting candidates even in ground based data is not difficult. They would be expected to be present in any star forming region containing significant numbers of stars with masses near $M_c$. For example, a recent CMD of an association in the Small Magellanic Cloud containing stars with ages $10-60$ Myr contains three stars whose colors and magnitudes are consistent with those of young white dwarfs (\\refto{dem}). Future HST observations of NGC 1818 and other young LMC clusters, to determine the cooling sequence, ages and masses of the white dwarf population in these clusters, should allow us to determine $M_c$ more precisely, and possibly for a range of metallicities. ", + "conclusions": "We have identified a candidate luminous white dwarf in the young star cluster NGC 1818 in the LMC. The object is $\\sim 35 $ arcsec from the cluster center (about $4.5 r_c$ and $0.6 r_h$). It has coordinates 5:04:13.8, $-$66:26:33.4 (J2000). In the HST passbands it has $V_{555}=18.43$, $(V_{555}-I_{814})=-0.26$ and $(U_{336}-V_{555})=-1.67$. In the Johnson-Cousins system this corresponds to $V=18.44$, $(V-I)=-0.25$ and $(U-V)=-1.32$. These values are corrected for reddening assuming $E(B-V)=0.05$. Posisson errors are $\\pm 0.03$ for $V$ and $\\pm 0.04$ for the colors. These do not include uncertainties in the transformation to the Johnson-Cousin system. The temperature is probably $\\gtsim 20,000$ K but is poorly constrained by the $(U-V)$ color. With the adopted distance modulus of 18.5, the object has absolute magnitude $M_V=-0.06$. If this object is indeed a white dwarf, then its mass is probably $ 1.1-1.3 \\msun$ (\\refto{Weid}). The composition of white dwarfs formed from high mass progenitors is expected to be Oxygen--Neon--Magnesium, but may be Carbon/Oxygen. A spectroscopic determination of its composition is a priority. If spectroscopic followup observations confirm the identity of the candidate star as a luminous young white dwarf, we have strongly constrained the critical mass at which stars stop evolving to type II supernovae to $M_c \\gtsim 7.6 \\msun$. \\smallskip \\noindent {\\bf Note:} A preliminary spectrum of the white dwarf candidate was obtained at the AAT 5 March 1998. The spectrum rules out the possibility that the object is a quasar. The velocity is indistinguishable from that of two other cluster members so it is also very unlikely to be a foreground object. Detailed modelling of the spectrum is currently in progress. \\medskip \\noindent {\\bf Acknowledgements:\\/} This research was supported in part by a PPARC rolling grant. SS acknowledges the support of the European Union through a Marie Curie Individual Fellowship. MBD gratefully acknowledges the support of the Royal Society through a URF. Funding for JH was provided by a grant from the Cambridge Commonwealth Trust, and from Trinity College. We would like to thank Brian Boyle, Matt Burleigh, Helen Johnston and Ray Stathakis for obtaining a spectrum of the candidate white dwarf. \\newcommand{\\journ}[4]{ {\\sl #1\\/} {\\bf #2}, #3 (#4).}" + }, + "9802/astro-ph9802321_arXiv.txt": { + "abstract": "In a previous paper we have written down equations describing steady-state, optically thin, advection-dominated accretion onto a Kerr black hole (\\cite{GP}, hereafter Paper I). In this paper we survey the numerical solutions to these equations. We find that the temperature and density of the gas in the inner part of the accretion flow depend strongly on the black hole spin parameter $a$. The rate of angular momentum accretion is also shown to depend on $a$; for $a$ greater than an equilibrium spin parameter $a_{eq}$ the black hole is de-spun by the accretion flow. We also investigate the dependence of the flow on the angular momentum transport efficiency $\\alpha$, the advected fraction of the dissipated energy $f$, and the adiabatic index $\\gamma$. We find solutions for $-1 < a < 1$, $10^{-4} \\le \\alpha \\le 0.44$, $0.01 \\le f \\le 1$, and $4/3 < \\gamma < 5/3$. For low values of $\\alpha$ and $f$ the inner part of the flow exhibits a pressure maximum and appears similar to equilibrium thick disk solutions. ", + "introduction": "In advection-dominated accretion flows (ADAFs) the accreting gas flows inward much more rapidly than it can cool. The energy released by accretion goes into heating the gas. This is in marked contrast to the usual thin accretion disk, where the radial velocity is small and the accretion energy is efficiently radiated away. If a black hole accretes via an ADAF much of the accretion energy can be carried across the horizon with the heated gas, reducing the luminosity well below that of a comparable thin disk. Paper I briefly summarizes the development of advection-dominated disk theory. Over the past few years, this theory has had notable success in reproducing the observed spectra of black hole candidate systems. For a review of the theory and applications of advection-dominated disk theory, see \\cite{n97}. Early models of advection-dominated flows around black holes (\\cite{ch96}, \\cite{nak96}, \\cite{nak97}, \\cite{cal97}, \\cite{nkh}) did not include a proper treatment of relativistic effects, but instead used an approximate pseudo-Newtonian potential due to \\cite{pw80}. Close to the black hole event horizon, the gas temperatures and velocities can become extremely high. This hot, rapidly rotating and infalling gas should produce important observable effects in the high-energy spectra of black hole candidates. Relativistic effects not accounted for by the pseudo-Newtonian potential are dominant in this innermost region. The character of the inner portion of the disk also depends strongly on the black hole spin $a$, which is not included in the pseudo-Newtonian treatment. For these reasons, it seems clear that a fully relativistic model is needed. Recently, advection-dominated disk models have begun to include the effects of general relativity, allowing a more accurate examination of the flow close to the event horizon. Both \\cite{acgl} (AGCL) and \\cite{pa97} (PA) have presented disk solutions in the Kerr metric. In Paper I, we wrote down a set of disk equations in the Kerr metric, and showed a few example solutions. Our relativistic disk equations differ in a number of respects from those of ACGL and PA. These differences are described in detail in Paper I. The main differences with ACGL are that we use a causal stress and include the relativistic enthalpy. PA also include these effects, although they use a more simplified prescription to enforce causality. Unlike PA, who use a polytropic equation of state, we solve the energy equation assuming that a constant fraction of the dissipated energy gets advected with the gas. Also, unlike ACGL and PA, we use the height prescription of \\cite{alp}. In this paper, we examine the structure of our numerical solutions for steady state flow onto a Kerr black hole in detail. In particular, we explore the effects of changing the dimensionless black hole spin $a$, the viscosity parameter $\\alpha$, the advected fraction of the dissipated energy $f$, and the adiabatic index $\\gamma_0$ (ACGL have shown solutions for three values of $a$, while PA presented a more extensive set of solutions for various values of $a$ and $\\alpha$). We show that there is an equilibrium spin rate for black holes accreting via ADAFs, above which the black hole will be spun down by accretion. We also show that there is a pressure maximum in the vicinity of the last stable orbit for low $\\alpha$ solutions (cf. NKH). Finally we show that the Bernoulli parameter can be positive in some regions of the flow, which at least suggests the possibility of a pressure-driven outflow. To familiarize readers with our notation, we summarize our relativistic disk equations and method of solution in \\S 2. In \\S 3 we examine the dependence of our disk solutions on $a$, $\\alpha$, $f$, and $\\gamma_0$. \\S 4 demonstrates that there is an equilibrium spin rate for black holes accreting from ADAFs. We discuss the implications of our results in \\S 5. ", + "conclusions": "\\subsection{Comparison to Earlier Solutions} Our disk equations differ in a number of respects from those used by other authors, as detailed in Paper I and above. Also, the solutions described above cover a wider range of parameter space than previous studies. We have found solutions over the ranges of $a=-0.99999$ to $0.99999$, $\\alpha = 0.0001-0.44$, $f = 0.01-1$, and $\\gamma = 1.3333-1.66$. Our sequence of solutions for $\\alpha = 0.001 - 0.3$ resembles the solutions of NKH in many respects. We find sonic radii ranging from $r_s \\simeq 4.28 - 9.41$ as $\\alpha$ increases, similar to the range $r_s = 4.20 - 10.63$ found by NKH. Our solutions also develop pressure maxima at small values of $\\alpha$, as discussed below. Since NKH used the pseudo-Newtonian potential, their solutions have some unphysical properties near the horizon, which are avoided by our use of the full relativistic equations. Our relativistic solutions with $a=0$ produce emission spectra which are qualitatively similar to those calculated from the NKH solutions. The densities tend to be somewhat higher near the horizon in our solutions due to the inclusion of relativistic effects, and this increases the luminosity of the flow substantially. This was shown by \\cite{n98}, who calculated the emission spectrum of Sgr A* using the relativistic dynamical solutions described in this paper. They were thus able to make a direct comparison between the spectrum computed from our solution and that computed from a pseudo-Newtonian solution as described by NKH. This showed that the inclusion of relativistic dynamics increased the emission by a factor of 10-100 in the infrared. Note that relativistic photon transport effects are not included in these spectra, apart from including the gravitational redshift. Also, the effects of using rotating black hole solutions have yet to be explored; however, it is clear that the substantial increases in temperature and density near the horizon in these solutions will have a dramatic effect on the emission spectrum. ACGL presented solutions with $a = 0$, 0.5, and 0.99 for $\\alpha = 0.1$, which included bremsstrahlung cooling. The shape of the $l(r)$ curves is similar to ours; however, the solutions differ in some other respects. First, the ACGL solutions have larger values of $l$ at the inner edge of the flow. They find $l_{in} \\simeq 3.2, 2.6, 1.7$ for $a = 0, 0.5,0.99$, respectively, whereas our solutions have $l_{in} = 2.14, 1.79, 0.93$ for these same values of $a$. One reason for this is that their $l(r)$ curves only extend in to the Schwarzschild radius at $r=2$, even for the $a=0.5$ and $a=0.99$ solutions, where the horizon sits at smaller radii. Also, in ACGL's scheme, the angular momentum eigenvalue which we call $j$ is equal to $l_{in}$ (or in their notation, $\\sL = \\sL_0$), whereas in our solutions, $j$ can differ substantially from $l_{in}$ due to viscous torques. Our values of $j$ for these solutions are 2.62, 2.31, and 1.73, which are somewhat closer to the ACGL values. Another difference between our solutions and those of ACGL is that their solutions show a maximum in $\\log c_s$ and inflections in $\\log P$ close to the inner edge, whereas our solutions have $\\log T$ and $\\log P$ increasing smoothly all the way in to the horizon. PA have presented the most extensive survey of solutions to date. They find two types of viscous disk solutions: ``type I'' solutions at low values of $\\alpha$ ($\\alpha = 0.001-0.045$ in the solutions shown) and ``type II'' solutions at high $\\alpha$ ($\\alpha \\geq 0.3$ in the solutions shown). One major difference between PA's scheme and ours is that they specify the angular momentum eigenvalue $j$ (which they call $L_0$) while we solve for it. Their type I solutions with $\\alpha$ ranging from 0.01 to 0.045 all have the same value of $L_0$, and show great variation in their angular momentum profiles at large $r$, whereas our solutions tend to have rather similar values of $l$ at large $r$. PA's type I solutions also have a maximum in the sound speed outside the sonic point, and the sound speed then decreases down to the horizon. This may be a consequence of their polytropic equation of state, which requires that the sound speed must decrease inward if the radial velocity increases inward more rapidly than $\\sim r^{-1}$. Their lowest-$\\alpha$ solutions at $\\alpha = 0.001 - 0.008$ have a super-Keplerian region in the inner part of the disk. In the type II solutions shown by PA, the sonic point occurs far out in the flow at $r \\sim 30-60$. This is a much larger sonic radius than in our high-$\\alpha$ solutions: our $\\alpha = 0.3$ solution has $r_s = 9.41$. According to PA, their type II solutions have the sonic point located at an outer critical point rather than an inner critical point as the Type I solutions do. Unlike PA, who found two types of solutions, and \\cite{ch96}, who has found solutions where the flow goes through radial shocks, we find smooth solutions for the entire range of parameter values, and we see no evidence for sudden transitions between different types of solutions. In order to look for alternate solutions, we varied $r_v$ and solved for the outer and middle sections of our solutions to see whether these sections would match up for an alternate value of $r_v$. Despite looking for alternate solutions with several choices of parameters, including low values of $\\alpha$, we found no additional solutions. It is worth commenting in some detail on why our solutions are shock-free while those of \\cite{ch96} are not. The origin of this difference lies in how the boundary conditions are treated. The difference is most easily explained by analogy with the problem of spherical (Bondi-Hoyle) accretion. To solve the Bondi problem one specifies the density and temperature, but not the radial velocity (nor, equivalently, the accretion rate), at large distance from the accreting object. The velocity at large radius is adjusted until the flow passes smoothly through the sonic point. This velocity, or the accretion rate, is thus an eigenvalue of the problem and must be solved for self-consistently. This approach has been validated by numerically solving a realistic initial value problem and showing that it settles down to the Bondi solution. If one were to specify the radial velocity or accretion rate the flow would not generally pass smoothly through the sonic point; rather, it would shock. Our treatment is analogous to the standard treatment of the Bondi problem, except that in our case the mass accretion rate is specified, while the angular momentum accretion rate $j$ is the eigenvalue. It is adjusted so that the flow passes smoothly through the sonic point. Chakrabarti's approach, on the other hand, is analogous to prespecifying the mass accretion rate in the Bondi problem: he fixes $j$. As a result the flow does not generally pass smoothly through the sonic point and his solutions contain one or more shocks. Our intuition, and the analogy with the Bondi problem, suggest that this is not the correct approach. The issue can only be settled conclusively, however, by solving a realistic initial value problem and showing that it converges to one solution or the other. This has not yet been done. \\subsection{Local Pressure Maxima} Our low-$\\alpha$ solutions have a maximum in pressure, density, and temperature in the inner disk. Figure 7 shows density contours in the inner parts of the flow for various values of $a$ and $\\alpha$. The contour plots are in the $x = r\\sin(\\theta), z = r\\cos(\\theta)$ plane, and assume that $\\rho \\propto \\exp(-\\theta^2/(2 H_\\theta^2))$. These solutions resemble the low-$\\alpha$ solutions of NKH. They are also similar to the thick disk models developed by a number of workers (see \\cite{fm76}, \\cite{pw80}, \\cite{rbbp}). In particular they have a pressure maximum, approximately constant specific angular momentum, and a region of super-Keplerian rotation where pressure support is important. Generally the pressure maximum lies close to the last stable orbit, and the sonic point lies close to the marginally bound orbit. Our approximations are most accurate for these low-$\\alpha$ solutions in that $H/R$ is relatively small (see Figure 2) and the low viscosity implies a disk with relatively low turbulent velocities. It is not entirely clear that the low-$\\alpha$ models are relevant, however. A lower limit to the efficiency of angular momentum transport is set by the existence of the global hydrodynamic instability of \\cite{pp84}. Simulations of the nonlinear development of the magnetorotational instability (\\cite{bgh95}) suggest an even larger lower limit of order $\\alpha \\sim 0.01$. For $\\alpha$ as large as this a pressure maximum does not develop, and the pressure increases smoothly down to the event horizon. A new result of this study is the similarity of low-$f$ solutions to low-$\\alpha$ solutions. Our $f=0.01$ and $f=0.03$ solutions show pronounced density and pressure maxima in the same region of the flow, just outside the last stable orbit. They also have super-Keplerian rotation in this region. However, the low-$f$ solutions have the sonic point at $r_s \\simeq 5.72$ for $f=0.03$ and at $r_s \\simeq 5.85$ for $f=0.01$, closer to the last stable orbit than the marginally bound orbit. \\subsection{Bernoulli Parameter} The Bernoulli parameter $Be$ measures the sum of the kinetic energy, potential energy, and enthalpy of the gas. Narayan \\& Yi (1994, 1995) pointed out that self-similar advection-dominated flows have a positive Bernoulli parameter for $f > 1/3$. The positivity of $Be$ suggests the possibility of a pressure-driven outflow; $Be$ is conserved for adiabatic, inviscid flows, so that gas with $Be > 0$ could flow outward adiabatically and still have positive kinetic energy at large radius. In its relativistic form, $Be = \\eta \\sE - 1$, where $\\sE \\equiv -u_t$. Figure 8 shows $Be$ for the solutions shown in Figs. 1--4, illustrating the variation of $Be$ with radius and with $a$, $\\alpha$, $f$, and $\\gamma$. Solutions with small values of $a$ reach only small positive values of $Be$. In these solutions, $Be$ tends to peak at $r \\sim 6-20$, and reaches peak values of $Be < 0.01$ (see Fig. 8b,d). Solutions with $a$ approaching unity can have substantially larger values of $Be$ which peak at the horizon (Fig. 8a). Solutions with $f < 1$, which radiate away a fraction $1-f$ of the dissipated energy, have $Be < 0$. At large radii, the solution with $f=0.3$ has $Be$ only slightly negative, as expected from the self-similar result that ADAFs with $f > 1/3$ have $Be > 0$. Variations in the Bernoulli parameter are directly related to the radial viscous energy flux, since the total radial flux of mass energy is conserved. We have \\begin{equation} Be = {\\dot{E}\\over{\\dot{M}}} - 1 + {t_t^r\\over{\\rho u^r}} \\end{equation} where $t_{\\mu\\nu}$ is the viscous stress tensor. Recall that $\\dot{E}$ at the event horizon is the actual rate of change of the black hole mass; furthermore, $\\dot{E} = const. \\simeq \\dot{M}$ when $f = 1$. One can show that, if $S$ is the shear stress measured in the local rest frame of the fluid, then \\begin{equation} \\label{ENERFLUX} t_t^r = -r\\sD u^\\phi S. \\end{equation} Since $u^\\phi \\sim \\sD^{-1}$ at the horizon, $t_t^r$ is finite at the horizon. Figure 9 shows the run of $-t_t^r$, the {\\it outward} viscous energy flux, for solutions with several values of $a$ but otherwise with the standard parameters. The existence of a finite outward angular momentum and energy flux at the horizon suggests a violation of causality, yet our solution is manifestly causal. How can this be? It turns out that these fluxes appear because of how the flow is divided into a mean and fluctuating part. Consider a simple example: an accretion flow consisting of a turbulent, unmagnetized fluid with negligible density fluctuations. At a given event, the flow has angular momentum $l + \\delta l$ and radial four-velocity $u^r + \\delta u^r$. The mean flow is defined so that $\\<\\delta l\\> = 0$ and $\\<\\delta u^r\\> = 0$, where the brackets denote an average over $t,\\theta,\\phi$. Then the outward flux of angular momentum is \\begin{equation} {T_\\phi}^r = (\\rho + u + p)(l + \\delta l)(u^r + \\delta u^r), \\end{equation} using the definition of the perfect fluid stress tensor. Averaging, \\begin{equation} {T_\\phi}^r = (\\rho + u + p)[l u^r + \\<\\delta l \\delta u^r\\>]. \\end{equation} The first term in brackets is due to the mean flow; the second term is what we have called the ``viscous'' angular momentum flux. Evidently the outward flux of angular momentum (or energy; the same considerations apply to ${T_t}^r$) appears because of how we have divided the flow into a mean and fluctuating part. Correlations in the fluctuations merely bias the angular momentum of accreting fluid elements. Causality is preserved. \\subsection{Assumptions and Limitations} Finally, it is worth offering a frank discussion of the assumptions behind our solutions and their limitations. A somewhat hidden assumption is that the accreting plasma is a two-temperature plasma with proton temperature much greater than the electron temperature. Only then will $f \\simeq 1$. This can be true only if most of the ``viscous'' dissipation goes into the protons and there is no collective effect that efficiently couples the protons and electrons. No such effect has yet been convincingly demonstrated to exist. The assumption that $f$ is constant with radius is only likely to hold true if $f \\simeq 1$ throughout the flow. We have nonetheless calculated solutions for constant $f < 1$ in order to illustrate the effects of a smaller value of $f$ on the dynamical aspects of the flow. In the future, it is clear that models will need to include cooling processes and calculate $f$ self-consistently. This task is complicated by the importance of Compton cooling in these flows, which depends not only on the local conditions at a particular radius, but also on the incident photon flux from all other radii. Thus far, cooling has been included in some detail in some models, but no model has included both a full treatment of cooling and fully relativistic dynamics. Some additional assumptions and limitations of our solutions are tied to our treatment of angular momentum transport. For example, we have vertically averaged the flow. This should produce a reasonably reliable solution close to the midplane, but is not predictive for flow near the poles. We could produce a full axisymmetric steady-state solution, but because it is not known how turbulent angular momentum transport varies with height in the accretion flow, such an effort would not significantly improve the reliability of the solution. In addition, the flow has been assumed smooth and steady. This is likely to be true only in a time-averaged sense. In particular, if $\\alpha \\sim 1$, then the turbulence that transports angular momentum will be only marginally subsonic and shocks and substantial density variations are likely." + }, + "9802/astro-ph9802098_arXiv.txt": { + "abstract": "We report the discovery of a fairly bright transient during observations with the Wide Field Cameras on board the BeppoSAX satellite in September 1996. It was detected at a peak intensity of 0.1 Crab (2 to 10 keV) and lasted between 6 and 40 days above a detection threshold of 2~mCrab. Two very bright type I X-ray bursts were detected from this transient in the same observations. These almost certainly identify this X-ray transient as a low-mass X-ray binary with a neutron star as compact component. The double-peaked time history of both bursts at high energies suggests a peak luminosity close to the Eddington limit. Assuming this to be true implies a distance to this object of 4~kpc. ", + "introduction": "\\label{secintro} Currently a program is carried out to regularly monitor the galactic bulge in 2 to 25 keV X-rays with the Wide Field Cameras (WFC) on board the BeppoSAX satellite. The main purpose of this program is to monitor relatively weak and short transient activity from various types of sources, in particular low-mass X-ray binaries (LMXBs). According to a recent count (Van Paradijs 1995), near to 30\\% of all $\\sim$130 known LMXBs are transient in nature. This population is concentrated in the sky towards the direction of the galactic center (e.g., Van Paradijs \\& White 1995, White \\& Van Paradijs 1996). We here present the discovery of a relatively bright transient during observations in September 1996, which exhibited X-ray bursts, and discuss its timing and spectral behavior in X-rays. In Sect.~\\ref{secobs}, we discuss the observations, in Sect.~\\ref{sectdetection} the detection and position of the transient, in Sect.~\\ref{secslow} trends in the intensity and spectrum, in Sect.~\\ref{secbursts} the two bright X-ray bursts that were detected and in Sect.~\\ref{secdisc} we evaluate the data. ", + "conclusions": "\\label{secdisc} Of the two types of X-ray bursts found in many X-ray binaries (see review by Lewin, Van Paradijs \\& Taam 1995), type I bursts are attributed to thermonuclear flashes on or near a neutron star surface. Detection of type I bursts is, therefore, a strong indicator for a neutron star. One diagnostic clearly distinguishes type I from the other type of bursts: only type I bursts exhibit spectral softening. A further characteristic of type I bursts is that they have black body spectra with temperatures up to a few keV. Thus, we can identify the two bursts reported here as type I bursts and conclude that there is strong evidence for the neutron star nature of this X-ray source. Type~I X-ray bursts have been seen from 42 galactic X-ray binaries according to Van Paradijs (1995). Although not all of them have confirmed optical counterparts, those 19 that do are all LMXBs. All 42 bursters have been classified as LMXBs, directly through the identification of the optical counterpart or indirectly through characteristics of the X-ray emission or association with a globular cluster. It is, therefore, almost certain that \\src\\ too is a LMXB with an as yet unidentified optical counterpart. The double-peaked nature of both bursts at high energies is indicative of near-Eddington luminosities (e.g., Lewin~\\ea 1995). This is supported by the black body temperatures which are similar to the (likewise high) values obtained for other bursts that reach the Eddington limit (e.g., Lewin at al. 1995). If interpreted as such, an estimate can be obtained of the distance. Assuming a 1.4~M$_{\\odot}$ neutron star with an Eddington limit of 2~10$^{38}$~erg~s$^{-1}$ and an observed peak bolometric flux of $(1.3\\pm0.3)~10^{-7}$~erg~s$^{-1}$cm$^{-2}$, the distance is 4~kpc. The galactic latitude of -8.1$^{\\rm o}$ makes a distance closer than the galactic center (8.5~kpc) indeed likely. Such a close distance suggests a reasonable perspective to find an optical counterpart and we urge optical observers to follow up on the position here published. The observed peak flux of the steady emission is about 2~10$^{-9}$~erg~s$^{-1}$cm$^{-2}$ in 2 to 10 keV which for the power law spectrum extrapolates to a 0.4 to 10~keV luminosity of ($6\\pm2$)~10$^{36}$~erg~s$^{-1}$ and for the thermal bremsstrahlung spectrum to $(5\\pm1)~10^{36}$~erg~s$^{-1}$. The uncertainty in these numbers is due to that in $N_{\\rm H}$, a distance of 4~kpc is assumed. The 0.4 to 10 keV peak luminosity is somewhat low though not unheard of within the group of LMXB transients (e.g., Chen~et~al. 1997). The distance of 4~kpc implies a burst emitting sphere radius of 8~km. This supports the neutron star identification. It is interesting to note that \\src\\ was not initially reported from data obtained with the all-sky monitor on board the Rossi X-ray Timing Explorer. The source was probably too near to the sensitivity of this instrument which for sources in uncrowded fields is about 50~mCrabs per dwell (3$\\sigma$, Levine \\ea 1996). The closest bright source is X1822-371 at 3.5~degrees. Therefore, it appears that the source was never much brighter than the peak intensity observed with BeppoSAX-WFC." + }, + "9802/astro-ph9802267_arXiv.txt": { + "abstract": "The evolution of optically selected quasars is usually supposed to be well described by a single constant evolution parameter, either $k_L$ or $k_D$, depending whether we refer to luminosity or density evolution. In this paper we present a study of the variations of the evolution parameters with redshift, for different cosmological models, in order to probe the differential evolution with redshift. Two different quasar samples have been analyzed, the AAT Boyle's et al. and the LBQS catalogues. Basically, these samples are divided in redshift intervals and in each of them $k_L$ and $k_D$ are estimated by forcing that $\\langle V/V_{max}\\rangle=0.5$. The dependence with respect to the cosmological parameters is small. Both AAT and LBQS show roughly the same tendencies. LBQS, however, shows strong fluctuations, whose origin is not statistical but rather due to the selection criteria. A discussion on selection techniques, biases and binning effects explains the differences between these results. We finally conclude that the evolution parameter is almost constant in the redshift range $0.7 \\leq z \\leq 1.7$, at least within $2\\sigma$, while it decreases slightly afterwards. Results depend on the binning chosen (but not in a very significant way). The method has been tested with Monte-Carlo simulated catalogues in order to give a better understanding of the results coming from the real catalogues. A correlation between $k_L$ ($k_D$) and $\\langle V/V_{max}\\rangle$ is also derived and is used for the calculation of the error bars on the evolution parameter. ", + "introduction": "\\label{intro} Since the first application of the $V/V_{max}$ test by Schmidt in 1968, the luminosity function of quasars is known to undergo a strong evolution, in the sense that the density of the most luminous quasars was far higher in the past. This was firstly interpreted in terms of either Pure Density (PDE - Schmidt, 1968) or Pure Luminosity Evolution (PLE - Mathez, 1976), and phenomenologically modeled with a single free parameter law. More complex models began to appear a few years later, such as a Luminosity Dependent Density Evolution (LDDE), proposed by Schmidt and Green in 1983. However, there wasn't any privileged model until the late 80's, when Boyle et al. (1988) favored PLE as a best fit of the observed luminosity function of a UVX sample. A similar recent result is found for X-Ray selected samples of quasars (Page et al. 1997; \\cite{Ja}; \\cite{Ba97}). The currently observed evolution of quasars appears however to be more complex nowadays, since the first indication of a reversing evolutionary trend around a redshift in the range [2.5,3] was observed (Shaver, 1994; Schmidt, Schneider and Gunn 1995; Warren, Hewett and Osmer 1994; Pei, 1995). A steepening of the luminosity function towards high redshifts is advocated by Goldschmidt and Miller (1997). Meanwhile, the evolution of the cosmic star formation density was found to be strikingly similar to that shown by Shaver for QSOs, although the maximum of the star formation density is attained at a lower redshift (Madau, 1996). In spite of this difference, the similarity of the variations of QSO luminosity density and of the field galaxy star formation has been interpreted as a clue that both phenomena could be closely linked (Boyle and Terlevitch, 1998; Silk and Rees 1998). Furthermore, the similarity is even more striking (both curves have a maximum around the redshift z=2.5) after applying the necessary correction for dust extinction in the density of high redshift galaxies (\\cite{Sa98}). This correction may be as high as a factor of ten since these galaxies are observed in their rest-frame ultraviolet. In the same time, complex models of this evolution begin to appear: $\\bullet $ At high redshift, the growth of Dark Matter halos according to Press-Schechter formalism and the parallel Eddington-limited growth of accreting Massive Black Holes (MBH) (\\cite{HL}) are likely to induce a decrease of density with increasing redshift ('negative DE') (Haiman and Loeb, 1997; Cavaliere and Vittorini, 1998; Krivitsky and Kontorovich, 1998; Novosyadlyj and Chornij, 1997). $\\bullet $ Around a redshift z=2.5, there could be two phenomena: a transition from high to low efficiency in advection-dominated flows, followed by the decline of accretion rate, giving luminosity proportional to $(1+z)^k$ with $k$ slightly variable around 3 (\\cite{Yi}). Alternatively, growing galaxies could assemble in groups of typically $5 \\; 10^{12} \\, M_{\\odot}$, where tidal effects refuel the MBH at lower and lower rate, translating into 'positive LE' (\\cite{CV}). $\\bullet $ At intermediate redshift, viscous instabilities induce long term, high amplitude variations of the accretion rate. The fraction of time spent at each luminosity level, convolved with the mass distribution gives the luminosity function and its evolution (\\cite{SE}). Galaxy collisions provide a mechanism which fuels galactic nuclei with gas in dense environments, giving raise to quasars in low luminosity galaxies (\\cite{LKM}). $\\bullet $ In addition, gravitational lensing mimics luminosity evolution in flux-limited samples, amplifying the more distant quasars, but whose detailed effect is not known, in particular the induced bias selection effects in magnitude limited quasar samples. Van Waebeke et al. (1996) define a new test, the $V/V_{max}$ statistics, and apply it to the AAT quasar sample (Boyle et al., 1990). They show that, contrary to the usual $\\langle V/V_{max}\\rangle $ test, the $V/V_{max}$ statistics leads to constraints on cosmological parameters. However these results rely on an arbitrary model of PLE, since the evolutionary effects dominate the cosmological effects in the QSO distribution, and we have to be sure of the reliability of the evolution model before constraining the cosmological parameters. So, the analysis in Van Waebeke et al. (1996) is rather a test of compatibility between an evolution model and a couple of cosmological parameters and there is a need for an {\\it independent} better understanding of the QSO evolution before applying such cosmological tests. A number of projects aim at assembling complete and homogeneous samples of several thousands of quasars which will allow more reliable and more subtle analyses, leading, thus, to verifications or improvement of complex theories on quasar evolution. Finally, there is some hope for progress in the understanding of the evolution of the stellar formation rate and of quasars, and in the determination of cosmological parameters compatible with the observed quasar distribution. Note that this is a necessary framework in order to make a precise estimate of the effects of the re-ionization in the future temperature maps of Planck-Surveyor. The aim of this paper is to explore in details the possibility for the evolution to depend on the redshift, and on whether we can distinguish between artefact and physical evolution. For the present study we decided to adopt a cosmology and an evolution model with a single evolution parameter which is allowed to depend on the redshift (and/or the magnitude in Section 4). This should lead to more complex evolution laws, which may help in theoretical understanding of quasar evolution and in turn could be the basis of future, more realistic cosmological tests. The technical basis of this work is the $\\langle V/V_{max}\\rangle $ test, performed in bins of redshift. In \\S \\ref{samples} and \\ref{evomodels} we make a brief description of the quasar samples and the evolution models used. In \\S \\ref{evomag} we give a justification of our choice to study quasar evolution in redshift bins and in \\S \\ref{m1} we present the different redshift binning modes used for our study. \\S \\ref{simulcat} contains the results obtained from the simulated catalogues. \\S \\ref{kvsz} contains the results for the real quasar samples. More precisely, it shows the measured dependence of the evolution parameter versus redshift, for a few sets of cosmological parameters. Finally, in \\S \\ref{discuss} we discuss the main results of our study. ", + "conclusions": "\\label{discuss} This paper presents an analysis of the evolution rate of quasars at various epochs. Our method differs from most previous ones in that it does not make use of any determination of luminosity function. It solely consists in computing the evolution parameters $k_L$ and $k_D$ (assuming PLE and PDE respectively) in bins of redshift, such that the $V/V_{max}$ are uniformly distributed. However, the various selection criteria applied in the construction of the two catalogues (AAT and LBQS) as well as the different redshift and magnitude limits make a direct comparison of the results quiet difficult. The method has been tested with Monte-Carlo simulated catalogues. Three binnings have been adopted, the first one with equal numbers per bin, the second and the third with a priori similar evolution rates inside each bin, according to the evolution law. Our results are the following: \\,\\, $\\bullet$ Both samples roughly show the same large trends, however modulations do appear in the LBQS results, which are likely to be correlated with the crossing of the main emission lines from the blue to the red side of the available spectral range. \\,\\, $\\bullet$ All results on the evolution parameters are quite similar whatever the hypothesis, PLE or PDE, mainly due to the inefficiency of the $V/V_{max}$ test to distinguishing between density and luminosity evolution. \\,\\, $\\bullet$ Similar results are also obtained for both power law and exponential parametrizations. Certainly a single phenomenon cannot be described by two different laws but we note that there is no significant difference between these laws in the redshift range $z \\in [0.7,2.2]$, as seen in the inset plot in Fig. \\ref{rellum}. The determination of quasars' evolution rate at high redshifts will also determine which of the proposed models is the appropriate one (if any). \\,\\, $\\bullet$ In bins with equally distributed evolution, however, both parameters $k_L$ and $k_D$ show far less variations with redshift and within error bars we can suppose that they are constant ($2.5 \\leq k_L \\leq 4, 1.5 \\leq k_D \\leq 3$ in a power law parametrization). \\,\\, $\\bullet$ $k_L$ ($k_D$) and $\\langle V/V_{max}\\rangle$ are linearly correlated, as shown in Fig. \\ref{correl1}. We confirm that a PLE (or PDE) with a constant evolution parameter is a good approximation for the redshift range $z \\in [0.7,1.7]$. At larger redshift, the results of our analysis of the LBQS are consistent with a maximum luminosity (or density) around $z=2.5$, but the reliability of these results is not yet established due to likely selection biases. The determination of this $k(z)$ dependence towards larger $z$ is essential for the understanding of the birth and growth of quasars at high redshifts, and the relation of the quasar phenomenon with star bursts in the primordial universe." + }, + "9802/astro-ph9802179_arXiv.txt": { + "abstract": "Following previous suggestions of other researchers, this paper discusses the prospects for astrometric observation of MACHO gravitational microlensing events. We derive the expected astrometric observables for a simple microlensing event with either a dark or self-luminous lens, and demonstrate that accurate astrometry can determine the lens mass, distance, and proper motion in a very general fashion. In particular we argue that in limited circumstances ground-based, narrow-angle differential astrometric techniques are sufficient to measure the lens mass directly, and other lens properties (distance, transverse motion) by applying an independent model for the source distance and motion. We investigate the sensitivity of differential astrometry in determining lens parameters by Monte Carlo methods, and derive a quasi-empirical relationship between astrometric accuracy and mass uncertainty. ", + "introduction": "In 1986 Paczy\\'{n}ski (\\cite{Paczynski86b}) suggested that photometric observations of gravitational microlensing might be used to indirectly study the population of massive compact objects in the galaxy, and in particular MAssive Compact Halo Objects (MACHOs) that might be a significant component of the dark matter thought to exist in the galaxy by dynamical considerations. Paczy\\'{n}ski's 1986 paper, and the observational proposals it fostered were met with some skepticism. However, the past several years have seen Paczy\\'{n}ski's suggestion spectacularly confirmed -- at the time of this writing four separate groups have reported significant numbers of candidate gravitational microlensing events from photometric observations of LMC, SMC, and galactic bulge sources. The large majority of the light curves for these candidate microlensing events match theoretical expectations for single lens objects, and all collaborations report a significant excess of microlensing event candidates above the number expected from known stellar populations. In particular, from their first two years of data the MACHO collaboration reports eight events toward the LMC where only one is expected from known stellar populations, and estimates that roughly half of the expected dark matter in the galactic halo is in the form of dark stellar mass objects (\\cite{Alcock96}). The difficulty in interpreting the MACHO collaboration events is that they are observed photometrically, which does not uniquely determine the mass of the lens -- instead the MACHO collaboration bases their conclusions on interpreting their event sample observables (namely event duration) in the context of a halo model (\\cite{Alcock96}). However, other interpretations are possible (see \\cite{Sahu94}, \\cite{Zhao97}, and \\cite{Gates97} for recent work arguing against a halo interpretation). Clearly it is desirable to measure MACHO physical properties in a model-free context. This objective has led a number of authors to propose the astrometric observation of MACHO gravitational microlensing events (\\cite{Hog95,Miyamoto95,Walker95}), a specialized application of an earlier suggestion by Hosokawa et al (\\cite{Hosokawa93}). In particular, Miyamoto and Yoshi proposed the separate astrometric observation of both lensing images in MACHO microlensing events (a small misuse of the term microlensing -- see \\cite{Paczynski86a}), and developed the theory of such astrometry. As we shall argue below, we find this suggestion implausible because of the small separation of the images. Instead, herein we consider astrometry of the lensed center-of-light, primarily for dark lenses. We find, as did Miyamoto and Yoshi, that high-precision astrometric observation of such microlensing events allows the estimation of the lens parameters (mass, distance, proper motion) appealing only to the properties of lensing. Moreover, we find that in a limited set of circumstances, the problem of determining a subset the lens parameters (mass, relative parallax, relative proper motion) is amenable to narrow-angle differential astrometric techniques very similar to those proposed and employed in gravitational companion search programs (\\cite{Shao92,Lestrade94,Benedict95,Gatewood96}). In particular we argue that if a suitable astrometric reference frame can be established, the lens mass can be directly measured by ground-based differential astrometric techniques independent of additional assumptions, and the lens distance and transverse velocity can be estimated by appealing to an independent model of source distance and proper motion (see similar remarks by \\cite{Walker95}). In such circumstances, many of the current issues regarding the nature of the lensing objects can be resolved. In this paper we assess the ability of astrometry to probe the physical parameters of microlensing events in which the lens is dark, with a particular emphasis on MACHO microlensing events. In \\S \\ref{sec:microlensing_theory} we introduce the theory to analyze a microlensing encounter as observed by a (near) terrestrial instrument, in terms particularly oriented toward narrow-angle differential astrometry. In \\S \\ref{sec:astrometric_observations} we address astrometric sensitivity to microlensing parameters through Monte Carlo techniques. Finally, in \\S \\ref{sec:discussion} we place our results in the context of envisioned astrometric instrumentation, discuss the near term prospects for such an astrometric program, and mention future extensions to this work. ", + "conclusions": "that high-precision astrometric observation of MACHO microlensing events hold the promise of determining fundamental lens parameters (mass, proper motion, transverse velocity) in a model-independent fashion. We maintain that it is sufficient to measure the motion of the center-of-light. In particular, in many cases accurate differential astrometry is sufficient obtain the lens mass without additional assumptions (in slight contrast to the earlier suggestion by \\cite{Walker95}, who ignored the value of parallactic effects), and reasonable lens distance and transverse velocity estimates can be obtained from an independent model of source distance and proper motion. Such narrow-angle differential astrometry is possible from the ground (\\cite{Shao92}). Alternatively, wide-angle microarcsecond-class astrometry can simultaneously determine the source and lens position and kinematic parameters without external information, and again the lens mass. Clearly the potential for probing the physical parameters of the putative halo object population by astrometric techniques is enormous. A program to probe microlensing events photometrically detected in the galactic bulge seems plausible for the planned Keck Interferometer (KI -- \\cite{Keck97}). KI requires a bright guide star to track atmospheric fluctuations of the interferometric fringes. The brightest bulge objects are 16th magnitude, within the fringe tracking capabilities for the two 10 m apertures. The expected 10 -- 20 $\\mu$as astrometric performance of KI yields microlensing parameter estimates sufficient to constrain lens mass and distance parameters for individual events, which will give profound insight into the nature of these objects. However, events in the LMC or SMC are not detectable from the Keck site, both because of geography and sensitivity. The brightest objects in the LMC are 17--18 magnitude, arguably fainter than the tracking capabilities of the KI. However, the declination of these fields would require KI zenith angles that severely degrade the astrometric performance. A large aperture astrometric interferometer in the Southern hemisphere such as the VLT interferometer (VLTI -- \\cite{Luthe94}) could measure Magellenic cloud events, and in particular determine the lens mass and distance with sufficient accuracy to resolve many of the current issues regarding their nature. Such measurements would be very challenging but are clearly very compelling. A number of authors have specifically suggested the application of planned space-based global astrometric techniques to analyze these events (\\cite{Hog95,Miyamoto95,Paczynski97}). In space-based applications astrometric references can be drawn from a global astrometric frame tied to extragalactic objects, and the positions, proper motions, and parallaxes of these references are known to a few microarcseconds in a quasi-static frame. Thus the necessity of establishing a narrow-angle relative frame for differential astrometry is removed. Further, astrometry at late times identifies the source proper motion and parallax in the global frame -- thereby establishing the source motion and distance. With the source distance and kinematics established, the lens parameters are all uniquely determined. Assuming ground-based observations measure lens masses and distance in the manner we describe, we believe the role of space-based microlensing astrometry programs by planned astrometric space missions such as SIM (\\cite{Unwin97}) and GAIA (\\cite{Lindegren96}) will be in probing the precision positions and particularly kinematics of the lensing objects. If ground-based measurements of Magelenic cloud events are not possible, then SIM and GAIA seem well-suited to offer definitive answers on the nature of the lenses. Finally, in the near future it is possible that CCD-based astrometry could make detections of microlensing astrometric perturbations, and possibly make rough estimates of lensing parameters, and/or breaking some of the degeneracies in photometric microlensing observations (see below). Pravdo and Shaklan (\\cite{Pravdo96}) report night-to-night astrometric repeatability of 200 $\\mu$as in data taken at the Palomar 5 m telescope, and speculate that limits might approach 100 $\\mu$as at the 10 m Keck Telescope. In further assessing these prospects we anxiously await the results of several nights of Keck observations recently made by Pravdo and Shaklan (\\cite{Shaklan97}). One of the key assumptions we have made in this work is the assumption of a dark lens. This assumption is plausible given the success photometric programs have had in fitting dark lens amplification models to photometric data. However, the instance of a luminous lens is possible and interesting (\\cite{Miralda96,Paczynski96b}), and the astrometric model derived here can be augmented in a straightforward way. Figure \\ref{fig:Llens_fitting} shows an example of a fit to a dataset generated with a luminous lens model unresolved from the lensed source. The physical parameters (bulge event, 10 kps source / 5 kpc lens distance, 0.1 M$_{\\sun}$ lens, lens/source brightness ratio of 0.1) were selected to be comparable to the example given in Figure \\ref{fig:model_fitting}. The model fit faithfully reconstructs the input physical parameters, including the correct attribution of source and lens brightness. One operational question that arises is how does one determine whether a dark lens model is appropriate for a given microlensing event. While one could straightforwardly test the luminous lens hypothesis by adding a relative source/lens luminosity parameter to the microlensing fit model as presented here, a more obvious and compelling resolution to this question is contained in the possible chromaticity of the astrometric observables as described in Eq.~\\ref{eq:astrometric_perturbation}. If the lens is luminous, then its spectral content is in general different from that of the source -- implying that both the photometric and astrometric observables will be functions of wavelength. Making the observations in a variety of spectral bands will identify the relative source-lens intensity and color, and provide the necessary data to robustly extend the microlensing model to the general case of luminous lenses. We defer a more systematic analysis of the luminous lens case to future work. \\begin{figure} \\epsscale{0.6} \\plottwo{Llens.photC.eps}{Llens.ast2C.eps} \\caption{Luminous Lens Microlensing Model Fitting. Here we show an instance of fitting a microlensing model with an unresolved luminous lens to synthetic terrestrial photometry and astrometry datasets for a microlensing encounter similar to that shown in Figure \\ref{fig:model_fitting}. The parameters for the event are a lens motion position angle of 30 deg, $p$ = 0.4, $r_E$ = 290 $\\mu$as, and $\\Pi$ = 100 $\\mu$as ($m$ = 0.1 M$_{\\sun}$) and a lens/source brightness ratio of 0.1. Again, we assume the event is identified photometrically, and differential astrometric measurements commence after that detection. The microlensing model described in \\S \\ref{sec:microlensing_theory} including a parameter for the relative lens/source brightness was simultaneously fit to both the photometric and astrometric data. Shown in each are the simulated data, true values, and the model fit. Left: the photometric lightcurve results. In this example we assume 3\\% photometry error. The time units on the $x$-axis are plotted relative to the barycentric $t_{max}$. Right: the corresponding depiction for the astrometry sequence relative to the nominal source position. The simulated 20 $\\mu$as astrometric measurements begin shortly before maximum magnification, and continue for 30 $t_{0}$ after maximum magnification. The fit is seen to faithfully reproduce the simulated datasets, and converge to the input model values including the appropriate source and lens brightness, even in the presence of frame drift. \\label{fig:Llens_fitting}} \\end{figure} The closely related problem of image blending is discussed in recent work by Alard (\\cite{Alard96}), and Wo\\'{z}niak \\& Paczy\\'{n}ski (\\cite{Wozniak97}), who consider the possibility that a second source (possibly unrelated to either source or lens) is unresolved from the image. Wo\\'{z}niak \\& Paczy\\'{n}ski find that the degeneracies in photometric observation of such events result in systematic errors in estimating lensing parameters. Based on our preliminary successes in correctly distinguishing source and lens luminosities, we concur with the speculation put forward by Wo\\'{z}niak \\& Paczy\\'{n}ski that multi-spectral astrometry and photometry breaks the degeneracy in (some subset of) blended events, and point this out as a particularly important case for future study. \\paragraph{Astrometric Detections -- Non-MACHO Events} A number of authors have suggested to broaden the applicability of the astrometric techniques to generic microlensing events (\\cite{Hosokawa93,Miralda96,Paczynski96b,Paczynski97}). These events could potentially be detected astrometrically in programs that concentrated on high proper motion objects (so as to sweep-out larger solid angles), or as a part of broader companion search program (something we have integrated into our PTI program -- \\cite{Colavita94}). While much of the phenomenology we have developed in \\S \\ref{sec:microlensing_theory} is directly applicable, there is a practical difficulty in establishing lensing parameters in such events by differential means. The first is the absence of a ready supply of reference objects that share common parallactic motions as the source. The fact that the rich LMC, SMC, and bulge fields used in the photometric surveys naturally yield an abundance of reference objects for which $\\delta D/D$ is small makes events in these fields unique. Without such a common parallactic reference, the systematic errors in the determination of $\\Pi$ will be too large to establish a precise lens mass from ground-based differential astrometry. Such events would seem to be best studied by space-based, global astrometric techniques. \\paragraph{Complex Lenses} While the majority of photometrically detected events are consistent with single lens hypothesis, a number of binary lens events have been reported (\\cite{Udalski94,Alard95a,Bennett96}). In a recent preprint Dominik (\\cite{Dominik97}) argues that photometry alone does not uniquely constrain the binary lens parameters. We speculate that additional astrometric information would break the degeneracies among various hypotheses in binary lens events through straightforward extensions of the astrometric theory developed here. We defer the analysis of the binary lens case to future work." + }, + "9802/astro-ph9802209_arXiv.txt": { + "abstract": "If primordial scalar and tensor perturbation spectra can be inferred from observations of the cosmic background radiation and large-scale structure, then one might hope to reconstruct a unique single-field inflaton potential capable of generating the observed spectra. In this paper we examine conditions under which such a potential can be reliably reconstructed. For it to be possible at all, the spectra must be well fit by a Taylor series expansion. A complete reconstruction requires a statistically-significant tensor mode to be measured in the microwave background. We find that the observational uncertainties dominate the theoretical error from use of the slow-roll approximation, and conclude that the reconstruction procedure will never insidiously lead to an irrelevant potential. \\vspace{24pt} \\noindent PACS numbers: 98.70.Vc, 98.80.Cq ", + "introduction": "Inflation produces metric perturbations, which are presently the most plausible cause of the observed temperature fluctuations in the cosmic background radiation (CBR) and which may act as the seeds for structure formation. During inflation the Hubble radius $H^{-1}$ ($H=\\dot{a}/a$ is the expansion rate) must have increased more slowly than the scale factor $a$, while during the radiation and matter eras it increased as $a^2$ and $a^{3/2}$, respectively. The dynamics of the evolution of the Hubble radius (or, equivalently, the evolution of the expansion rate $H$) during inflation is usually modelled by assuming that the energy density is dominated by scalar-field potential energy. The number of dynamical degrees of freedom associated with the evolution of the scalar-field potential energy (and hence the Hubble radius) is unknown. If there is only one relevant degree of freedom, that degree of freedom is called the {\\it inflaton,} and the scalar-field potential is called the {\\it inflaton potential.} Inflation models where a single inflaton field slowly evolves under the influence of the inflaton potential are called single-field slow-roll models. Such models are very attractive because of their simplicity, because they arise naturally in a host of particle physics models, and because other more complicated models can often be expressed in terms of an effective single-field slow-roll model. The assumption that only a single scalar field is dynamically relevant can, at least in principle, be tested by a series of consistency relations which represent relations between the scalar and tensor perturbations produced in single-field models \\cite{Recon,LLKCBA}. Inflation models where more than one scalar field is dynamically important are phenomenologically much more complicated, and almost certainly must be studied on a case-by-case basis. In this paper we shall restrict our discussion to the single-field case, where the scalar potential $V(\\phi)$ is the only piece of information to be specified. Reconstruction of the inflaton potential (see Ref.\\ \\cite{LLKCBA} for a review) refers to the process of using observational data, especially microwave background anisotropies, to determine the inflaton potential capable of generating the perturbation spectra inferred from observations \\cite{Recon}. Of course there is no way to prove that the reconstructed inflaton potential was the agent responsible for generating the perturbations. What can be hoped for is that one can determine a {\\it unique} (within observational errors) inflaton potential capable of producing the observed perturbation spectra. The reconstructed inflaton potential may well be the first concrete piece of information to be obtained about physics at scales close to the Planck scale. As is well known, inflation produces both scalar and tensor perturbations, which each generate microwave anisotropies (see Ref.\\ \\cite{LLrep} for a review). As reviewed in the next section, if $V(\\phi)$ is given the perturbation spectra can be computed exactly in linear perturbation theory through integration of the relevant mode equations. If the scalar field is rolling sufficiently slowly, the solutions to the mode equations may be approximated using the slow-roll expansion \\cite{SB,LL,LPB}. The standard reconstruction program makes use of the slow-roll expansion, taking advantage of a calculation of the perturbation spectra by Stewart and Lyth \\cite{SL}, which gives the next-order correction to the usual lowest-order slow-roll results. The biggest hurdle for successful reconstruction is that many inflation models predict a tensor perturbation amplitude well below the expected threshold for detection, and as we will see, even above detection threshold the errors can be considerable. If the tensor modes cannot be identified a unique reconstruction is impossible, as the scalar perturbations are governed not only by $V(\\phi)$, but by the first derivative of $V(\\phi)$ as well. Knowledge of only the scalar perturbations leaves an undetermined integration constant in the non-linear system of reconstruction equations. Another problem is that simple potentials usually lead to a nearly exact power-law scalar spectrum with the spectral index close to unity. In such a scenario only a very limited amount of information could be obtained about high energy physics from astrophysical observations. Recently, Wang et al.\\ \\cite{WMS} suggested an alternative danger to reconstruction. They studied the possibility that the potential might be sufficiently complicated that the slow-roll expansion breaks down, and a full integration of the mode functions is necessary to calculate the perturbation spectra to the accuracy expected by the next round of satellite observations by {\\sc map} and Planck. It is, of course, clear that any perturbative procedure such as the slow-roll expansion has some limit to its validity. But here, the breakdown of the slow-roll expansion is particularly troublesome; firstly because of the technical difficulty that there is no known method of computing yet higher-order corrections, and secondly because there is a general possibility that for sufficiently complicated potentials the perturbation series may not converge at all. However, the key concern for reconstruction is whether such a situation will lead to a {\\it misconstruction} of the inflaton potential. In this paper, we shall argue that for the purpose of reconstruction the breakdown of the slow-roll approximation is in many senses desirable. First of all, we show that there is no practical danger that one can be misled into thinking perturbative reconstruction is working when in fact it is not. Secondly, the very failure of the perturbative approach implies that extra parameters are required in order to explain the observed data set, and hence there is the prospect of obtaining a greater amount of information about the inflationary model. Finally, we will demonstrate that the theoretical errors from the use of slow-roll equations are sub-dominant to the expected observational errors from a satellite experiment such as Planck. Wang et al.\\ \\cite{WMS} described a particular, admittedly rather {\\em ad hoc}, scalar field potential, for which they demonstrated that the result of exact mode equation integration differed significantly, by the standards of say the Planck satellite, from the slow-roll expansion. We shall largely concentrate on that potential as a means of concrete illustration of the very general points we shall make. ", + "conclusions": "\\setcounter{equation}{0} We have found that a complicated potential, such as the example used by Wang et al.\\ \\cite{WMS}, would be a boon to reconstruction of the inflaton potential, as it would provide extra information accessible to observations, in the form of scale-dependence of the density perturbations. Such potentials are also more likely to have tensor modes at a detectable level, which is required for a complete reconstruction to be performed. We have analyzed the impact of observational errors on the reconstruction, and conclude that any theoretical errors from use of the slow-roll equations are likely to be sub-dominant. One might even be able to test this be solving the mode equation for the reconstructed potential, though we have not tried to pursue that route here. Finally, it has always been clear that there must be some limits to the applicability of perturbative reconstruction, although it does a good job even for the Wang et al.\\ potential. However, such a failure should be immediately evident from the data, as one would not be able to successfully match observations with a truncated Taylor series expansion of the power spectra. It would appear, therefore, that there is no danger that perturbative reconstruction might appear to be working, but in fact be producing an irrelevant potential." + }, + "9802/astro-ph9802273_arXiv.txt": { + "abstract": "We have used spectra taken between 1992 and 1997 to derive the spectroscopic orbit of the eclipsing double-lined spectroscopic binary HD~197770. This binary has a period of $99.69 \\pm 0.02$~days and K amplitudes of $31.2 \\pm 0.8$ and $47.1 \\pm 0.4$~km~s$^{-1}$ for components A \\& B, respectively. The $m\\sin^3i$ values for A \\& B are 2.9 and 1.9, respectively, and are close to the actual masses due to the eclipsing nature of this binary. Both components of HD~197770 have spectral types near B2~III. This means both components are undermassive by about a factor of five and, thus, evolved stars. Additional evidence of the evolved nature of HD~197770 is found in 25, 60, and 100~\\micron\\ IRAS images of HD~197770. These images show 2 apparent shells centered on HD~197770; a bright 60~\\micron\\ shell with a $14\\arcmin$ diameter and a larger ($1\\fdg 2$ diameter) bubble-like feature. At least one of the components of HD~197770 is likely to be a post-AGB star. ", + "introduction": "Interest in the star HD~197770 (HR~7940, $\\alpha(2000) = 20^{\\rm h}43^{\\rm m}13\\fs 52$, $\\delta(2000) = +57\\arcdeg 6\\arcmin 50\\farcs 9$) increased greatly with the discovery that its line-of-sight has a polarization feature coincident with the 2175~\\AA\\ extinction bump (\\cite{cla92}; \\cite{and96}). Out of the 30 sightlines with UV spectropolarimetry, such a polarization bump has only been seen along one other sightline (\\cite{wol97}). HD~197770 has long been known to have a variable radial velocity (\\cite{ada24}; \\cite{pla30}). Observations originally intended to study the sightline towards HD~197770 have shown it to be a double-lined spectroscopic binary (\\cite{han94}; \\cite{cla96}). Recent photometric observations have shown that HD~197770 is an eclipsing binary (\\cite{jer93}; \\cite{cla96}). We obtained spectra of HD~197770 between 1992 and 1997 in order to determine the spectroscopic orbit of HD~197770. The spectroscopic orbit, coupled with the eclipsing nature of the binary, allowed us to determine the masses of the binary components. ", + "conclusions": "Since HD~197770 exhibits shallow eclipses ($\\delta V \\sim 0.05$), its inclination must be near $90\\arcdeg$ (\\cite{cla96}). Thus, the values quoted in Table~\\ref{table_orbit} for $m\\sin^3 i$ are close to the actual masses of the components. Comparing an unpublished Pine Bluff Observatory (PBO) spectrum of HD~197770 to the spectra presented in Walborn \\& Fitzpatrick (1990), we find the spectral class to be B1~V-III or B2~III for both stars combined. Considering both the PBO spectrum and the $T_{\\rm eff}$ value determined in \\S\\ref{sec_analysis}, the most likely spectral type is B2~III. All lines in the spectra of HD~197770 are double, leading to the conclusion that both stars have similar spectral types. Assuming the radii of a B2~III star (12~R$_{\\sun}$; \\cite{dri97}) for both components, the inclination of this eclipsing binary is $\\geq 73\\arcdeg$. The masses of components A \\& B are then $\\leq 3.3$ and 2.2~M$_{\\sun}$, respectively. A normal B2~III star has a mass around 15 M$_{\\sun}$ (\\cite{dri97}). Thus, both components of this binary are undermassive for their given spectral types. This marks both components as evolved stars. \\begin{figure}[tbp] \\begin{center} \\plotone{all_four.eps} \\caption{The IRAS $2\\arcdeg \\times 2\\arcdeg$ images centered on HD~197770 are displayed. The image intensity ranges are 1.2-2.5~MJy~sr$^{-1}$ for the 12~\\micron\\ image (upper left), 4.2-4.95~MJy~sr$^{-1}$ for the 25~\\micron\\ image (upper right), 4.0-6.0~MJy~sr$^{-1}$ for the 60~\\micron\\ image (lower left), and 25-35~MJy~sr$^{-1}$ for the 100~\\micron\\ image (lower right). \\label{fig_all_four}} \\end{center} \\end{figure} HD~197770 lies in the Cygnus region on the edge of Cyg OB7 and Cep OB2. It seems to be very near and possibly on the edge of a molecular cloud/star formation region including Lynds~1036 and 1049. HD~197770 is associated with two IRAS point sources, IRAS~20420+5655 and 20418+5700, but it is clearly non-stellar on the IRAS map at 60~\\micron. Figure~\\ref{fig_all_four} displays the $2\\arcdeg \\times 2\\arcdeg$ region centered on HD~197770 for the IRAS 12, 25, 60, 100~\\micron\\ bands. The IR SED peaks at 60~\\micron\\ implying warm dust in the immediate vicinity of the binary (\\cite{gau93}). In addition to the 60~\\micron\\ shell with a radius of 14$\\arcmin$, the IRAS map shows the signature of an apparent bubble of cleared dust with a radius of approximately $0\\fdg 6$. This bubble is easiest to see in the 60 and 100~\\micron\\ images, but is also present in the 25~\\micron\\ image. The dark regions in the upper and lower right hand corners of the images mark the edge of a molecular cloud which has been mapped in CO (\\cite{dob94}). The molecular cloud wraps around three sides of HD~197770 roughly surrounding the evacuated area in the IRAS map. In particular, to the west of HD~197770 is a buried young stellar object (Lynds~1036) and to the east of HD~197770 is a pulsar (\\cite{dew85}). So sequential star formation seems possible although the pulsar may be a background object. In any event, the HD~197770 binary appears to have formed on the edge of this cloud and cleared out an area around it. The association of HD~197770 with this cloud allows us to use the distance to the cloud (440~pc; \\cite{dob94}) as an estimate of the distance to HD~197770. The existence of IR flux peaking at 60~\\micron\\ at the position of HD~197770 and bubble of cleared dust surrounding HD~197770 imply that this binary has undergone at least two episodes of mass loss. A measurement of the luminosity of this binary would greatly help in determining its evolutionary stage, but the distance to the system is not known. HIPPARCOS gives a parallax of $0.52 \\pm 0.50$~mas which results in a $3\\sigma$ lower limit on the distance of 495~pc. This is consistent with the ``guilt by association'' distance given in the last paragraph. By assuming the distance to HD~197770 is 440~pc, $L_B/L_A = 2$, and $T_{\\rm eff} = 21,000$~K, we can estimate the luminosity and radii of the components of this binary. The resulting luminosities and radii are $9.3 \\times 10^3$~L$_{\\sun}$ ($M_V = -3.2$) and 7.4~R$_{\\sun}$ for HD~197770~A and $4.7 \\times 10^3$~L$_{\\sun}$ ($M_V = -2.5$) and 5.2~R$_{\\sun}$ for HD~197770~B. These radii result in an inclination for the binary of $\\geq 81\\arcdeg$ and, thus, masses of $\\sim$3.0 and 2.0~M$_{\\sun}$ for HD~197770~A \\& B, respectively. We note the similarity between this system and the eclipsing binary $\\upsilon$~Sgr ($m_p\\sin^3 i = 2.5$~M$_{\\sun}$, $m_s\\sin^3 i = 4.0$~M$_{\\sun}$, \\& $a\\sin i = 210$~R$_{\\sun}$; \\cite{dud90}) for which the primary is a hydrogen-deficient A-type supergiant and the secondary has a spectral type of B2~Ib (\\cite{sch83}). According to Plavec (1973) and Schoenberner \\& Drilling (1983), the system is the result of type BB binary evolution, where the primary is now a post-AGB star. Uomoto (1986) has discussed the possibility that such systems are progenitors of Ib-type supernovae. Two kinds of additional observations are needed of this binary. First, high resolution (R~$\\geq 50,000$) and high S/N optical spectra are needed spaced throughout the $\\sim$100 day orbit. These observations would confirm the K-amplitude of component A which is currently only based on three measurements. Second, imaging with an optical interferometer would give an accurate distance measurement when coupled to the results of this work. The Navy Prototype Optical Interferometer could perform this imaging when it is fully operational (\\cite{arm98}). Calculation of each star's luminosity would help in understanding their evolution stage." + }, + "9802/astro-ph9802045_arXiv.txt": { + "abstract": "Two long-slit spectra of the diffuse ionized gas in NGC 891 are presented. The first reveals variations parallel to the major axis in emission line ratios in the halo gas at $z=700$ pc. It is found that filaments of H$\\alpha$ emission show lower values of [N$\\,$II]/H$\\alpha$, [S$\\,$II]/H$\\alpha$ and [O$\\,$I]/H$\\alpha$. Although this result is expected if the filaments represent the walls of evacuated chimneys, it merely reflects a more general correlation of these ratios with H$\\alpha$ surface brightness along the slit, and may simply arise from radiation dilution effects. Halo regions showing low line ratios are probably relatively close to ionizing sources in the disk below. The results highlight difficulties inherent in observations of edge-on galaxies caused by lack of knowledge of structure in the in-plane directions. The [S$\\,$II]/[N$\\,$II] ratio shows almost no dependence on distance along the major axis or H$\\alpha$ surface brightness. Values of [O$\\,$I]/H$\\alpha$ indicate that H is 80--95\\% ionized (assuming $T=10^4$ K), with the higher ionization fractions correlating with higher surface brightness. Much more interesting information on the nature of this gaseous halo comes from the second observation, which shows the vertical dependence of [N$\\,$II]/H$\\alpha$, [S$\\,$II]/H$\\alpha$, [O$\\,$I]/H$\\alpha$, and [O$\\,$III]/H$\\beta$ through the brightest region of the DIG halo. The most surprising result, in complete contradiction to models in which the DIG is ionized by massive stars in the disk, is that [O$\\,$III]/H$\\beta$ rises with height above the plane for $z>1$ kpc (even as [N$\\,$II]/H$\\alpha$, [S$\\,$II]/H$\\alpha$, and [O$\\,$I]/H$\\alpha$ are rising, in line with expectations from such models). The run of [S$\\,$II]/[N$\\,$II] is also problematic, showing essentially no contrast with $z$. The [O$\\,$III] emission probably arises from shocks, turbulent mixing layers, or some other secondary source of ionization. Composite models in which the line emission comes from a mix of photo-ionized gas and shocks or turbulent mixing layers are considered in diagnostic diagrams, with the result that many aspects of the data can be explained. Problems with the run of [S$\\,$II]/[N$\\,$II] still remain, however. There is a reasonably large parameter space allowed for the second component. For the photo-ionized component, only matter-bounded models succeed, putting a fairly strong restriction on the clumpiness of the halo gas. Given the many uncertainties, the composite models can do little more than demonstrate the feasibility of these processes as secondary sources of energy input. A fairly robust result, however, is that the fraction of H$\\alpha$ emission arising from the second component probably increases with $z$. From values of [O$\\,$I]/H$\\alpha$, H is essentially 100\\% ionized at $z=0$ kpc and 90\\% ionized at $z=1$ kpc (again assuming $T=10^4$ K). ", + "introduction": "The vast majority of the free electrons in the ISM of the Milky Way reside in a thick ($\\sim$ 900-pc scale height) diffuse layer known as the Reynolds layer or the Warm Ionized Medium (e.g. Reynolds 1993). This phase fills about 20\\% of the ISM volume, with a local midplane density of about 0.1 cm$^{-3}$. Such a phase is now known to be a general feature of external star-forming galaxies, both spirals (e.g. Walterbos 1997; Rand 1996) and irregulars (e.g. Hunter \\& Gallagher 1990; Martin 1997, hereafter M97), where it is commonly referred to as Diffuse Ionized Gas (DIG). However, for edge-on spirals, only in the more actively star-forming galaxies does the gas manifest itself as a smooth, widespread layer of emission detectable {\\it above} the HII region layer (Rand 1996). One such galaxy, NGC 891, is an attractive target for study, not only because of its prominent DIG layer (Rand, Kulkarni, \\& Hester 1990; Dettmar 1990), but also its proximity ($D=9.5$ Mpc will be assumed here) and nearly fully edge-on aspect ($i>88\\arcdeg$; Swaters 1994). One of the outstanding problems in the astrophysics of the ISM is the ionization of these layers. For the Reynolds layer, the local ionization requirement ($5\\times 10^6$ s$^{-1}$ per cm$^{2}$ of Galactic disk; Reynolds 1992) is comfortably exceeded (by a factor of 6 or 7) only by the ionizing output of massive stars. Alternatively, the ionization would require essentially all the power put out by supernovae (Reynolds 1984) -- hence, this energy source could contribute at some level but probably cannot explain all of the diffuse emission. Photo-ionization models, on the other hand, must explain how the ionizing photons can travel $\\sim$ 1 kpc or more from their origin in the thin disk of massive stars to maintain this distended layer. Crucial information on both the ionization and thermal balance of DIG comes from emission line ratios. In the Reynolds layer, ratios of [S$\\,$II] $\\lambda\\lambda6716,6731$ and [N$\\,$II] $\\lambda\\lambda6548,6583$ to H$\\alpha$ are generally enhanced relative to their HII-region values, while [O$\\,$III] $\\lambda5007$/H$\\alpha$ is much weaker. These contrasts are in accordance with models in which photons leak out of HII regions and ionize a larger volume, with the radiation field becoming increasingly diluted with distance from the HII region [Mathis 1986; Domg$\\ddot {\\rm o}$rgen, \\& Mathis 1994; Sokolowski 1994 (hereafter S94; see also Bland-Hawthorn, Freeman, \\& Quinn 1997)]. The effect of this dilution, measured by the ionization parameter, $U$, is primarily to allow species such as S and O, which are predominantly doubly ionized in HII regions, to recombine into a singly ionized state. The effect may be less noticeable for N because it is mostly singly ionized in HII regions. The Wisconsin H$\\alpha$ Mapper (WHAM) has been used to determine [O$\\,$I]/H$\\alpha$ in three low-latitude directions, resulting in values $<0.01$ to 0.04 (Haffner \\& Reynolds 1997). Such low values imply, since the ionization of O and H are strongly coupled by a charge exchange reaction, that the diffuse gas is nearly completely ionized (Reynolds 1989). Although weak, [O$\\,$III] emission has been detected in two directions in the Reynolds layer at $b=0\\arcdeg$ (Reynolds 1985), with the result [O$\\,$III]/H$\\alpha = 0.06$. Reynolds postulated that the [O$\\,$III] emission does not arise from diluted stellar ionization but from gas at about 10$^5$ K, presumably the same gas as seen in C$\\,$ IV $\\lambda$1550 and O$\\,$III] $\\lambda1663$ emission by Martin \\& Bowyer (1990). The origin of this rapidly cooling gas is unclear. [O$\\,$III] emission from the DIG of NGC 891 and the implications for DIG ionization is one of the main subjects of this paper. In external spiral galaxies, smooth increases in [S$\\,$II]/H$\\alpha$ and [N$\\,$II]/H$\\alpha$ vs. distance from HII regions have been observed in both the in-plane and vertical directions in accordance with photo-ionization models [Walterbos \\& Braun 1994; Dettmar \\& Schulz 1992; Rand 1997a (hereafter R97); Golla, Dettmar, \\& Domg$\\ddot {\\rm o}$rgen 1996; Greenawalt, Walterbos, \\& Braun 1997; Wang, Heckman, \\& Lehnert 1997]. The same trends are seen in irregulars (Hunter \\& Gallagher 1990; M97). This behavior has been revealed in NGC 891 through spectra using long slits oriented vertically to the plane. Dettmar \\& Schulz (1992) placed a slit at $R=65''$ NE of the nucleus, while R97 took a deeper spectrum at $R=100''$ NE. R97 found that [N$\\,$II]/H$\\alpha$ rises to a value of 1.4, implying a very hard ionizing spectrum. S94, which pays particular attention to modeling the DIG of NGC 891, can predict such a high value only by assuming a stellar IMF extending to 120 M$_{\\sun}$, a reduction in cooling efficiency due to elemental depletions, and hardening of the radiation field by the intervening gas. [O$\\,$I]/H$\\alpha$ was not detected by Dettmar \\& Schulz (1992) in the halo of NGC 891 at $R=65''$ NE, with an upper limit of 0.05. Dettmar (1992) also reported an upper limit on [O$\\,$III]/H$\\beta$ of 0.4 at the same location. A wealth of forbidden-line long-slit data on bright DIG and HII regions in irregular galaxies has recently been published by M97. Through the use of line-diagnostic diagrams (e.g. Baldwin, Phillips, \\& Terlevich 1981; Veilleux \\& Osterbrock 1987), she finds that while photo-ionization models can explain the line ratio behavior in many galaxies, the rather shallow fall-off of [O$\\,$III]/H$\\beta$ with distance from HII regions and the sharp rise in [O$\\,$I]/H$\\alpha$ seen in some galaxies imply a second source of ionization. Shocks are favored as the most likely second source. The forbidden lines, though bright, are sensitive to metallicity and temperature and thus their interpretation in terms of ionization scenarios is complicated by uncertainties in abundances, degree of depletion, and sources of non-ionization heating. A more direct constraint on the ionizing spectrum has come from the very weak He$\\,$I $\\lambda$5876 line. He$\\,$I/H$\\alpha$ is relatively easy to interpret in terms of the ratio of helium- to hydrogen-ionizing photons, allowing the hardness of the ionizing spectrum, the mean spectral type of the responsible stars, and the upper IMF cutoff to be inferred, assuming pure stellar photoionization. The results for the Reynolds layer (Reynolds \\& Tufte 1995), for HI worms from equivalent radio recombination lines (Heiles et al. 1996) and for NGC 891 (R97) all imply a much softer spectrum than do the forbidden lines. Further consequences of this discrepancy are discussed in the above three references. The goal of this paper is to make further progress in understanding the ionization of DIG in spirals. The motivations are two-fold. First, the DIG halo of NGC 891 features several bright filaments and shells. It is likely that some of these are chimney walls (Norman \\& Ikeuchi 1989) surrounding regions of space evacuated by many supernovae. In this case, radiation from any continuing star formation near the base of the chimney will have an unimpeded journey to the walls, and thus the filaments may be directly ionized and show a spectrum more like an HII region than diffuse gas, which receives a significant contribution from relatively soft diffuse re-radiation (Norman 1991). If true, then the filaments should show lower [N$\\,$II]/H$\\alpha$ and [S$\\,$II]/H$\\alpha$ than the surrounding gas. To this end, a spectrum has been taken with a slit oriented parallel to the major axis, but offset into the halo gas, traversing several filaments. The second purpose is to study in more detail the dependence of line ratios on $z$ beyond the results reported in R97 for He$\\,$I/H$\\alpha$ and [N$\\,$II]/H$\\alpha$. By adding measurements of [S$\\,$II], [O$\\,$III], [O$\\,$I], and H$\\beta$, one can form diagnostic diagrams and thus constrain the source(s) of ionization in the spirit of M97. ", + "conclusions": "Using a slit oriented parallel to, and offset 700 pc above, the major axis of NGC 891, a spectrum of the DIG has been taken which reveals a clear correlation of ratios of forbidden lines to H$\\alpha$ with the H$\\alpha$ surface brightness. The original motivation for this observation was to search for variations in line ratios on and off the filaments of DIG as further evidence that they are walls around evacuated chimneys. But although the filaments do show reduced [S$\\,$II]/H$\\alpha$, [N$\\,$II]/H$\\alpha$, and [O$\\,$I]/H$\\alpha$ relative to gas on adjacent lines of sight, the contrast merely reflects the overall correlation. The relationship probably indicates that regions of brighter H$\\alpha$ emission receive a radiation field with a higher ionization parameter. Also, although some of the filaments show deviations from the observed smooth trend of mean velocity with position along the major axis, these departures could simply be due to geometric effects: if the filaments are inner galaxy features, they will bias the mean velocity for their line of sight away from the systemic velocity. [S$\\,$II]/[N$\\,$II] surprisingly shows no significant variation along the slit. Finally, the H$\\alpha$-emitting halo gas at this height is about 80--95\\% ionized, based on the observed range of [O$\\,$I]/H$\\alpha$ and assuming $T=10^4$ K. The correlation of [O$\\,$I]/H$\\alpha$ with surface brightness probably reflects a higher degree of ionization where the photon field is more intense. Results from this observation emphasize the difficulty in interpreting DIG observations of edge-on galaxies. Confusion is caused by uncertainties in the location of a parcel of gas along the line of sight, its effective distance from a source of ionization and other unrelated gas in the same direction. It is difficult from such observations to draw conclusions about the environment of the filament, for example. Spectra from a slit oriented perpendicular to the plane at $R=100''$ along the major axis on the NE side show a rise of [S$\\,$II]/H$\\alpha$, [N$\\,$II]/H$\\alpha$, and [O$\\,$I]/H$\\alpha$ with $z$. At the midplane, [O$\\,$I]/H$\\alpha$ values indicate that H is essentially 100\\% ionized, dropping to 90\\% at $z=1$ kpc, assuming $T=10^4\\,$K. The $z$-dependence of these line ratios is expected if the gas is ionized by massive stars in the disk. However, it is unexpectedly found that [O$\\,$III]/H$\\beta$ also rises with $z$, whereas it should decline with $z$ in photo-ionization models. This result necessitates the consideration of secondary sources of ionization. [S$\\,$II]/[N$\\,$II] unexpectedly shows essentially no dependence on $z$ and H$\\alpha$ surface brightness. Put another way, [N$\\,$II]/H$\\alpha$ shows the same disk-halo contrast as that of [S$\\,$II]/H$\\alpha$, whereas a smaller contrast is expected. Strong [O$\\,$III] emission is expected from several energetic procesess. We considered shocks and turbulent mixing layers as sources of such gas. Models in which a small fraction of the H$\\alpha$ emission comes from one of these mechanisms can be made to fit the data reasonably well, but most noticeably the remarkable constancy of [S$\\,$II]/[N$\\,$II] with $z$ is still difficult to reproduce. In the case of shocks, it is difficult to constrain the shock speed or the contributed fraction of the DIG emission at this point. Of course, the line of sight may sample shocks with a range of speeds and some mean value. There is also significant latitude in the parameters in the case of TMLs. Other sources of strong [O$\\,$III] emission may include cooling galactic fountain gas and microflares from magnetic reconnection. These should be explored further in light of the current results. Because of these facts and possibilities, the results are meant only to indicate the feasibility of such classes of models and the likelihood that one or more physical processes is producing intermediate temperature gas in the halo of NGC 891. On the other hand, the finding that the second source of line emission becomes more important as $z$ increases may be reasonable. In the case of TMLs, for example, Shull \\& Slavin (1994) point out that this process may indeed be more common at large $z$, where superbubbles break out of the thin disk gas layers, producing Rayleigh-Taylor instabilities and shear flows that lead to the mixing. These authors were attempting to explain the larger scale-height of C$\\,$IV UV absorption line gas relative to N$\\,$V (Sembach \\& Savage 1992) as an increasing predominance of TMLs over SN bubbles with height off the plane -- the former producing higher C$\\,$IV/N$\\,$V. If the TML process begins only at the approximate height where breakout occurs, while the stellar radiation field is increasingly diluted with $z$, then an increasing fraction of H$\\alpha$ emission from TMLs may be quite reasonable. The rough fractions found in \\S3.2.3 are comparable to those expected by Slavin et al. (1993) for the Milky Way diffuse H$\\alpha$ emission. If the photo-ionized and secondary components of the DIG emission both arise in exponential layers with different scale-heights, then the composite models, although illustrative, can be used to estimate roughly the relative scale-heights. In the two shock models and one TML model considered, the fraction of emission arising from the second component is 3--5 times higher at $z=2$ kpc than at $z=0$ kpc. Assuming exponential layers, the scale-height of the second component must be 3--4 times that of the photo-ionized component. This conclusion is very tentative, however, given the uncertainties in the modeling and the lack of information on [O$\\,$III]/H$\\beta$ at higher $z$. It it is tempting to identify the second component with the high-$z$ tail of H$\\alpha$ emission found in the deeper spectrum of Rand (1997a). However, the exponential scale-height of this tail is 5--7 times that of the main component, and it contributes about 50\\% of the emission at $z=2$ kpc. While these numbers do not quite match those for the second component proposed here, there may yet prove to be a connection. For the photo-ionized component, we found the most success by using the model from S94 with the lowest terminal hydrogen column considered for individual clouds in the DIG layer. Using larger columns tends to push the model curves towards those of the radiation-bounded case, which is found to be very difficult to incorporate in a successful composite model. This constraint suggests that the DIG consists of quite small clumps (or filaments or sheets, since S94's calculation is one-dimensional) of several pc thickness, for a representative density of $n_e=0.1$ cm$^{-3}$. If this conclusion is not borne out by future observations, the composite models presented here will need to be reconsidered. The S94 model used here also features the hardest emergent stellar spectrum considered (in order to produce high [S$\\,$II]/H$\\alpha$ and [N$\\,$II]/H$\\alpha$), but more modest spectra may be allowable if other sources of non-ionization heating are at work (e.g. Minter \\& Balser 1997). Despite complications introduced by the second DIG component, it should be noted that the observed properties of the three ratios [S$\\,$II]/H$\\alpha$, [N$\\,$II]/H$\\alpha$, and [O$\\,$I]/H$\\alpha$ with $z$ are still reasonably explained {\\it to first order} by photo-ionization models alone. Their smooth increase with $z$ is as predicted as are their rough values. Emission from the second component probably has only a secondary effect on these ratios. The $z$-independence of [S$\\,$II]/[N$\\,$II] is not understood in either a pure photo-ionization or composite model. An undesirable aspect of the composite models considered here is that emission from photo-ionized and (for example) shock ionized gas is simply added together with no unified physical picture in mind. It would be desirable eventually to have, say, a calculation of the evolution of a superbubble which included photo-ionization, shocks and TMLs in a more self-consistent way. For instance, what is the effect of the radiative precursor of a shock which enters gas already ionized by dilute stellar radiation? The emission line properties revealed by the perpendicular slit share some similarities with those of the halo of the starburst galaxy M82. In the Fabry-Perot data of Shopbell \\& Bland-Hawthorn (1997), [N$\\,$II]/H$\\alpha$ shows a general tendency to rise with $z$ on the N side, up to the limit of measurability at about $z$=750 pc. On the S side, there is little dependence on z, with perhaps 0.6 typical. [O$\\,$III]/H$\\alpha$=0.03 at $z=0$ pc and 0.08 at $z=750$ pc (values of [O$\\,$III]/H$\\beta$ are about 3 times higher assuming little extinction). In a long-slit spectrum through the halo of M82, M97 sees higher values of [O$\\,$III]/H$\\beta$, reaching 0.7 at $z=1$ kpc. The sequence of points in her diagnostic diagram of [S$\\,$II]/H$\\alpha$ vs. [O$\\,$III]/H$\\beta$ has a rising slope, as in NGC 891, although much steeper (other irregulars show a falling or nearly flat slope). [O$\\,$I]/H$\\alpha$ also rises with $z$ in M82, and shows a range of values (about 0.01 to 0.1) similar to those reported here. Shopbell \\& Bland-Hawthorn (1997) point out that their ratios become more shock-like with distance from the starburst, as also noted by Heckman, Armus, \\& Miley (1990). This behavior is now seen in a spiral halo as well. Other evidence for multi-phase halos is provided by the study of NGC 4631 by Martin \\& Kern (1998). They detect an extensive halo of [O$\\,$III] emission which spatially coexists with the observed soft X-ray and H$\\alpha$ emitting halos. Within this halo are several bright [O$\\,$III] condensations in which the measured [O$\\,$III]/H$\\beta$ ratio is $\\sim 20$. On this basis, they argue that the H$\\alpha$ and [O$\\,$III] emission is tracing distinct components in a multi-phase halo medium. It should be pointed out that [O$\\,$III]/H$\\alpha$ values of 0.12--0.21 (comparable to values in NGC 891) are seen in the DIG of M31 (Greenawalt et al. 1997). There is little correlation of this ratio with the brightness of the DIG or its remoteness from visible HII regions. The high [O$\\,$III]/H$\\alpha$ in this case may be due to a radiation field not as dilute as in the halo of NGC 891 (for instance, [N$\\,$II]/H$\\alpha$ and [S$\\,$II]/H$\\alpha$ are substantially lower than in the halo of NGC 891 at $z=1-2$ kpc). Greenawalt et al. (1997) also find that TMLs may contribute some fraction of the H$\\alpha$ emission in regions of very faint DIG, but no more than about 20\\% and most likely only a few percent, similar to the findings for NGC 891. More light has been shed on the question of [O$\\,$III]/H$\\beta$ trends in the DIG of face-ons by Wang et al. (1997). For three of their five galaxies with good [O$\\,$III] detections in their DIG spectra, they find that [O$\\,$III]/H$\\beta$ is in the range $<0.2$ to 2 and, for a given slit, is systematically higher than in the HII regions in that slit. They also find that the [O$\\,$III] line widths are usually larger than those of [N$\\,$II], and thus refer to a \\lq\\lq quiescent\\rq\\rq\\ photo-ionized DIG component which accounts for the bulk of the H$\\alpha$, [N$\\,$II] and [S$\\,$II] emission, and a \\lq\\lq disturbed\\rq\\rq\\ component (shocks and TMLs are considered) contributing a minority ($<$ 20\\%) of the H$\\alpha$ emission but responsible for the [O$\\,$III] emission. For vertical hydrostatic equilibrium, the contrast in line-widths indicates that the scale-height of the disturbed component is 1.5--2 times greater than that of the quiescent DIG. In NGC 891, the line-widths are dominated by galactic rotation and thus such an analysis cannot be carried out. The existence of a second source of line emission may be relevant for the issue of the low He$\\,$I/H$\\alpha$ ratio. The value of 0.027 for the lower halo of NGC 891 implies a much softer spectrum than is required to explain the high [N$\\,$II]/H$\\alpha$ and [S$\\,$II]/H$\\alpha$ (R97). Apart from the possibility that the forbidden line emission is complicated by additional sources of ionization and non-ionizing heating, He$\\,$I/H$\\alpha$ may also be affected by secondary ionization sources if they contribute sufficient H$\\alpha$ emission. For instance, the ratio is fairly sensitive to shock conditions. The SM 100 km s$^{-1}$ model gives a value of 0.027. The ratio is 0.005 for an 80 km s$^{-1}$ shock, but this model predicts insignificant [O$\\,$III] emission. The Dopita \\& Sutherland (1996) higher velocity, lower density (n=1), magnetized models give a much more significant ionized precursor. For the lowest velocity considered, 150 km s$^{-1}$, the shock itself gives a low ratio of 0.019, but the line fluxes for the precursor are not given. For a 200 km s$^{-1}$ shock, the ratio from combined shock and precursor is 0.046. Slavin, Shull, \\& Begelman (1993) do not predict He$\\,$I emission. Regardless, if shocks or TMLs can provide, say 25\\% of the H$\\alpha$ emission at $z=2$ kpc, then there may be a region of parameter space which can produce low enough He$\\,$I/H$\\alpha$ so that the composite line ratio is significantly reduced below that of the stellar-ionized gas alone. The contributions of these sources required in \\S 3 may not be sufficiently large, but the effect is worth future consideration. The inferred stellar temperature, mean spectral type and upper IMF cutoff would then all be underestimated. Finally, it is worth re-emphasizing that all emission line fluxes and ratios presented here are averaged along a line of sight through the DIG layer, and that local variations in, for example, the derived ionization fraction of H surely exist. Also, although the vertical dependence of the line ratios has been very revealing, only one slit position has been observed, covering the halo above the most active region of star formation in the disk, and close to an H$\\alpha$ filament. A key question is how these halo properties vary with environment. Does [O$\\,$III]/H$\\beta$ show the same behavior above more quiescent parts of the disk? Is this behavior peculiar to the halo of NGC 891 only, or is it a general feature of DIG halos? These questions will be addressed by further observations. \\vspace*{0.5in} The author has benefited from many useful discussions about DIG ionization from R. Reynolds, R. Walterbos (whose comments as referee also improved the paper), J. Slavin, R. Benjamin, J. Shields, and others. The help of the KPNO staff is also greatly appreciated." + }, + "9802/astro-ph9802335_arXiv.txt": { + "abstract": "We report on results of recent, high resolution hydrodynamic simulations \\index{simulations, hydrodynamic}% of the formation and evolution of X-ray clusters \\index{X-ray clusters}% of galaxies carried out within a cosmological framework. We employ the highly accurate piecewise parabolic method (PPM) \\index{piecewise parabolic method}% on fixed and adaptive meshes which allow us to resolve the flow field in the intracluster gas. The excellent shock capturing and low numerical viscosity of PPM represent a substantial advance over previous studies using SPH. We find that in flat, hierarchical cosmological models, the ICM \\index{ICM}% is in a turbulent state long after turbulence \\index{turbulence}% \\index{ICM!turbulence}% generated by the last major merger should have decayed away. Turbulent velocites are found to vary slowly with cluster radius, being $\\sim 25\\%$ of $\\sigma_{vir}$ in the core, increasing to $\\sim 60\\%$ at the virial radius. We argue that more frequent minor mergers maintain the high level of turbulence found in the core where dynamical times are short. Turbulent pressure support is thus significant throughout the cluster, and results in a somewhat cooler cluster ($T/T_{vir} \\sim .8$) for its mass. Some implications of cluster turbulence are discussed. ", + "introduction": "Our conception of galaxy clusters\\footnote{to appear in {\\em Ringberg Workshop on M87}, eds. K. Meisenheimer \\& H.-J. R\\\"{o}ser, Springer Lecture Notes in Physics, 1998.} as being dynamically relaxed systems has undergone substantial revision in recent years. Optical observations reveal substructure \\index{clusters!substructure}% in 30-40\\% of rich clusters (Geller \\& Beers 1982; Dressler \\& Shectman 1988). A wealth of new X-ray observations have bolstered these findings, providing evidence of recent mergers in clusters previously thought to be archtypal relaxed clusters (e.g., Briel \\etal ~1991). Also eroding the conventional view has been the success of ``bottom-up\" or hierarchical models \\index{hierarchical models}% of cosmological structure formation in accounting for the formation of galaxies and large scale structure \\index{galaxies!formation}% \\index{galaxies!large scale structure}% \\index{galaxies!X-ray clusters}% in the universe (e.g., Ostriker 1993). Within such models, a cluster sized object is built up through a sequence of mergers of lower-mass systems (galaxies $\\rightarrow$ groups $\\rightarrow$ clusters). In a flat universe ($\\Omega_o = 1$) as predicted by inflation, \\index{inflation}% mergers would be ongoing at the present epoch. In open models ($\\Omega_o < 1$), mergers cease at a redshift $z \\sim \\Omega_o^{-1}-1$, and clusters become relaxed by today. The amount of substructure observed in X-ray clusters of galaxies at $z \\sim 0$ is thus a powerful probe of cosmology. Evrard \\etal ~(1993) and Mohr \\etal ~(1995) have explored this ``morphology-cosmology\" connection, \\index{morphology-cosmology connection}% and concluded that a high $\\Omega$ universe is favored. Interestingly, Tsai \\& Buote (1996) reach the opposite conclusion. Cluster mergers \\index{cluster mergers}% \\index{simulations!cluster mergers}% have been explored numerically by several groups (Schindler \\& M\\\"{u}ller 1993; Roettiger, Loken \\& Burns 1997; Roettiger, Stone \\& Mushotzky 1998). In these hydro/N-body simulations, two hydrostatic King models are collided varying the cluster--subcluster mass ratio. It is found that major mergers induce temperature inhomogeneities and bulk motions in the ICM of a substantial fraction \\index{ICM, bulk motions}% of the virial velocity ($> 1000 ~km/s$). Roettiger \\etal ~suggest that these bulk motions may be responsible for the observed temperature substructure seen in some X-ray clusters, as well as bending Wide-Angle Tailed radio sources, \\index{WAT radio sources}% energizing cluster radio halos, \\index{clusters!radio halos}% and disrupting cooling flows. \\index{clusters!cooling flows}% If hierarchical models are correct, the thermal and dynamical state of the ICM could be considerably more complex than the above mentioned simulations indicate. In a flat universe, for example, the ICM would be constantly bombarded by a rain of minor mergers in addition to the occasional major merger. Also omitted in those simulations are a variety of cosmological effects which may be important, including memory of the complex formation history of the merging clusters, infall of matter along filaments, accretion shocks, large scale tides, and cosmic expansion. In this paper we show results of numerical simulations that take all these effects into account. We find in two flat models investigated, that quite generally the ICM of rich galaxy clusters is in a turbulent state. The turbulent velocities are typically 60\\% the virial velocity at the virial radius, decreasing inward to roughly 25\\% within the core. The relatively slow decline in turbulence amplitude with decreasing radius suggests that frequent minor mergers are an important driving mechanism in addition to rare massive mergers. In addition, we find ordered fluid circulation in the core of one well--resolved cluster which is likely the remnant of a slightly off-axis recent merger. ", + "conclusions": "We have shown using high resolution hydrodynamic simulations that the ICMs in bright X-ray clusters in flat hierarchical models are turbulent throughout. The turbulence in strongest in the outskirts of the cluster and weaker in the core. Due to the declining temperature profile in cluster halos, the turbulence is found to be mildly supersonic ($M \\sim 1.6$) near $r_{vir}$, decreases rapidly to $M \\sim 0.5$ at $\\sim \\frac{1}{3}r_{vir}$, and thereafter declines more slowly to $M \\sim .3$ in the core. Here we argue that infrequent major mergers cannot sustain the observed level of turbulence in the core. It is known from simulations of decaying turbulence in a box that the turbulent kinetic energy decays as $t^{-\\eta}$ where t is measured in units of the dynamical time. \\index{turbulence!decay rates}% The exponent $\\eta$ depends weakly on the nature of the turbulence, but is around 1.2 for compressible, adiabatic, hydrodynamic turbulence (Mac Low \\etal ~1998). The time for a sound wave to propagate from the center of the cluster SB to a radius $.01, .1, 1 \\times r_{vir}$ is $.014, .173, 3.1 ~Gyr$, respectively. The cluster underwent a major merger at $z=0.4$, or $5.2 ~Gyr$ earlier. Taking the sound crossing time as the dynamical time, we predict that fluid turbulence induced by the major merger at $z=0.4$ would have decayed to $.006, .017, .56$ of its initial value by $z=0$. Several possibilities suggest themselves to account for the high fluid velocity dispersions seen in the core. The first is that energy is somehow pumped into the core by motions in the outer parts of the cluster which relax on longer timescales. However, shock waves generated by supersonic motions in the outskirts would weaken into acoustic disturbances as they propagated into the dense, hotter core. Gravitational accelerations in the core would be dominated by the local dark matter distribution which would relax on a timescale comparable to the turbulence decay timescale. Another pumping mechanism discussed by Roettiger, Burns \\& Loken (1996) is global oscillations of the cluster potential following a major merger. They find that rms velocities decay to $\\sim 200 ~km/s$ by 2 Gyr after core passage, and remain quite constant thereafter. This is substantially less than the velocities we find. The second possibility, which we consider more likely, is that core turbulence is driven by the more frequent minor mergers. Lacey \\& Cole (1993) have quantified the merger rates in hierarchical models. They find that the merger rate for CDM scales as $(\\Delta M/M_{cl})^{-\\frac{1}{2}}$ where $\\Delta M$ is the subcluster mass. Whereas most of a cluster's final mass is typically accreted in a single major merger, they find that the cluster will typically accrete $\\sim 10\\%$ of its mass in ten minor mergers of clumps $\\sim 1\\% $ of its final mass. The most probable formation epoch for a $10^{15} M_{\\odot}$ cluster in the stardard CDM model we have simulated is at .7 $t_{Hubble}$, or 4 Gyr ago. The mean time between minor mergers is thus 0.4 Gyr---comparable to the dynamical time at a tenth the virial radius. Is there sufficient energy in minor mergers to sustain the turbulence in the core, and if so, how is the energy deposited? The kinetic energy of ten $10^{13} M_{\\odot}$ subclusters is $\\sim 10^{63}$ erg, as compared to approximately $10^{62}$ erg of turbulent kinetic energy within $0.1 r_{vir}$. Thus, a 10 \\% energy conversion efficiency is required for this mechanism to be correct. If the coupling is purely hydrodynamic (i.e., shocks), then the energy available is the kinetic energy of the gas in the subcluster, which is down by a factor of $\\Omega_b$ from the estimate above. Since $\\Omega_b =.05 $, this energy is insufficient. Thus, it would seem that a substantial gravitational coupling between the ICM and the dark matter in the subclusters is required. This is equivalent to saying that the gas remains bound to the subcluster until it reaches the core. Roettiger \\etal ~(1996) found that this is indeed the case. There are a number of interesting implications to significant levels of turbulence in the cores of X-ray clusters, many of which have already been pointed out by Roettiger \\etal ~(1996), including Doppler shifting of X-ray emission lines, bending of Wide-Angle Tailed radio galaxies, and powering cluster radio halos. Our findings strengthen their conclusions. For example, the turbulent amplification of magnetic fields would be expected to be most efficient in cluster cores where dynamical timescales are shortest. Moreover, continuous stirring by minor mergers could modify cooling flows appreciably. Because turbulent pressure ``cools\" inefficiently compared to atomic processes, turbulent pressure support could become increasingly important in the central parts of a cooling flow. Its effect would be to reduce the mass inflow rate into the cluster center. Secondly, at radii much less than the cooling radius, turbulent motions would concentrate cooling gas into filaments, and possibly account for the observed $H \\alpha$ filaments. Finally, we note that ordered circulation in the cores of X-ray clusters such as we have found might account for the S-shaped symmetry of radio tails seen in some sources (e.g., M87; B\\\"{o}hringer \\etal ~(1995), Owen, these proceedings.) \\index{M87!radio tails}% {\\it Acknowledgements:} This work was partially supported by grants NASA NAGW-3152 and NSF ASC-9318185. Simulations were carried out on the Connection Machine-5 and Silicon Graphics Power Challenge Array at the National Center for Supercomputing Applications, University of Illinois." + }, + "9802/astro-ph9802103_arXiv.txt": { + "abstract": "We describe a model-independent method of assessing the uncertainties in cross-correlation lags determined from AGN light curves, and use this method to investigate the reality of lags between UV and optical continuum variations in well-studied AGNs. Our results confirm the existence of such lags in NGC 7469. We find that the continuum variations at 1825\\,\\AA, 4845\\,\\AA, and 6962\\,\\AA\\ follow those at 1315\\,\\AA\\ by $0.22^{+0.12}_{-0.13}$ days, $1.25^{+0.48}_{-0.35}$ days, and $1.84^{+0.93}_{-0.94}$ days, respectively, based on the centroids of the cross-correlation functions; the error intervals quoted correspond to 68\\% confidence levels, and each of these lags is greater than zero at no less than 97\\% confidence. We do not find statistically significant interband continuum lags in NGC 5548, NGC 3783, or Fairall 9. Wavelength-dependent continuum lags may be marginally detected in the case of NGC 4151. However, on the basis of theoretical considerations, wavelength-dependent continuum lags in sources other than NGC 7469 are not expected to have been detectable in previous experiments. We also confirm the existence of a statistically significant lag between X-ray and UV continuum variations in the blazar PKS\\,2155$-$304. ", + "introduction": "Over the last ten years, a number of intensive monitoring experiments on active galactic nuclei (AGNs) have been carried out (for reviews, see Netzer \\& Peterson 1997 and Peterson 1993). While the primary purpose of many spectroscopic monitoring campaigns has been to determine the response of the broad emission lines to continuum variations and thus determine the structure and kinematics of the line-emitting gas through the process of ``reverberation mapping'' (Blandford \\& McKee 1982), these campaigns have also provided an opportunity to search for time delays between different continuum bands. Such time delays are expected to exist if the continuum flux in one waveband is reprocessed continuum emission from another waveband, either directly through irradiation or indirectly through viscous processes of the emitting plasma. The standard accretion-disk model for AGNs may exhibit both. Recently, Wanders et al.\\ (1997) found evidence for such time delays in the UV spectra of the Seyfert 1 galaxy NGC 7469 obtained with the {\\it International Ultraviolet Explorer (IUE)}\\/ spacecraft. It was found that the continuum flux variations around 1700--1800\\,\\AA\\ lag behind those at 1315\\,\\AA\\ by $0.3\\pm0.07$\\,d. This result was subsequently supported by contemporaneous optical spectra (Collier et al.\\ 1998) that showed that the continuum variations around 4825\\,\\AA\\ lag behind the 1315\\,\\AA\\ variations by $1.2\\pm0.3$\\,d, and those around 6925\\,\\AA\\ lag by $1.7\\pm0.7$\\,d. The increase in time lag with wavelength in NGC\\,7469 is strong evidence for continuum reprocessing models of AGNs in which the longer-wavelength photons are reprocessed shorter-wavelength photons originating closer to the central source. A legitimate question to ask at this point is why such an effect has not been previously reported in the literature, even though other AGNs have been monitored in a similar fashion? This question is of particular interest in the case of the well-studied galaxy NGC 4151, which was monitored at an even higher sampling rate than NGC 7469 during a 10-day campaign in 1993 (Crenshaw et al.\\ 1996; Kaspi et al.\\ 1996; Warwick et al.\\ 1996; Edelson et al.\\ 1996). Other Seyfert galaxies that have been well-monitored simultaneously in the UV and optical are NGC 5548 (Clavel et al.\\ 1991; Peterson et al.\\ 1991, 1992; Korista et al.\\ 1995), NGC\\,3783 (Reichert et al.\\ 1994; Stirpe et al.\\ 1994), and Fairall\\,9 (Rodr\\'{\\i}guez-Pascual et al.\\ 1997; Santos-Lle\\'{o} et al.\\ 1997). None of these studies found a significant lag between different continuum waveband flux variations. Why were such lags found for NGC\\,7469, but not for the others? The answer to this question is not trivial. The detection of such lags is dependent on the sampling characteristics of the light curves, as well as on the variations of the light curves themselves, i.e., the auto-correlation function (ACF) of the light curve. In any event, the formal statistical significance of a time-lag detection through a cross-correlation analysis is hard to assess. There is no generally agreed upon way to estimate the errors in cross-correlation lags and thus attach a level of significance to a time-lag detection. Indeed, interband continuum lags {\\it have} been reported in previous campaigns, but the detections were not thought to be statistically significant. In this paper, we will introduce what we believe is a conservative model-independent approach to estimating the uncertainties in cross-correlation lag determinations (\\S\\,2) and use it to re-examine existing data sets (\\S\\,3). We will then discuss the implications of these results (\\S\\,4) and present our conclusions (\\S\\,5). ", + "conclusions": "In this paper, we have introduced a model-independent Monte-Carlo method for estimating conservatively uncertainties in cross-correlation lags, and have tested this method for a variety of conditions. We have applied this method to the UV/optical light curves of several well-studied non-blazar AGNs and the blazar PKS\\,2155$-$304. The reality of wavelength-dependent continuum lags found in NGC 7469 (Wanders et al.\\ 1997; Collier et al.\\ 1998) is supported by this analysis at at level of confidence higher than 95\\%. We confirm at about a 95\\% confidence level the detection of a $\\sim$2-hr X-ray/UV lag in PKS\\,2155$-$304. However, we find no evidence for statistically significant UV/optical continuum lags in the non-blazar AGNs NGC 5548, NGC 3783, and Fairall 9, and we find only marginal evidence for the effect in NGC 4151. By scaling other AGNs relative to NGC 7469, we find however that the absence of significant wavelength-dependent continuum lags in other AGNs does not appear to be inconsistent with a thermal accretion-disk model --- most AGNs have not been monitored at high enough temporal frequency for UV/optical continuum lags to be detectable. Failure to detect such lags in previous experiments does not pose a serious threat to accretion-disk models, as long as the variability signal is propagated at the speed of light. We are grateful for support of this study by NASA (through ADP grant NAG5-3497 and LTSA grant NAG5-3233) and by the National Science Foundation (through grant AST94-20080). The authors thank P.T.\\ O'Brien and an anonymous referee for a number of suggestions. \\clearpage" + }, + "9802/astro-ph9802225_arXiv.txt": { + "abstract": "We present optical data which shows that G\", the optical counterpart of the $\\gamma$-ray pulsar Geminga, pulses in B with a period of 0.237 seconds. The similarity between the optical pulse shape and the $\\gamma$-ray light curve indicates that a large fraction of the optical emission is non-thermal in origin - contrary to recent suggestions based upon the total optical flux. The derived magnitude of the pulsed emission is $m_B = 26.0 \\pm 0.4$. Whilst it is not possible to give an accurate figure for the pulsed fraction (due to variations in the sky background) we can give an upper limit of $m_B \\approx 27 $ for the unpulsed fraction. ", + "introduction": "The nature of the bright $\\gamma$-ray source Geminga remained elusive from the first observations using SAS-B (\\cite{fich75}) until its recognition as a pulsar with a period of 0.237 seconds in $\\gamma$ rays (\\cite{bert92} \\cite{big92}) and in X-rays (\\cite{hal92}). Based upon colour considerations an optical candidate was proposed, G\" with a m$_V$ of 25.5 (\\cite{halp88}). This star had a measurable proper motion (\\cite{big93}) indicating a probable distance of about 100 pc and thereby making a probable association with a neutron star. Subsequent Hubble Space Telescope observations have given a distance based upon parallax of $159^{+59}_{-34}$ pc (\\cite{car96}). Optical observations in B showed Geminga to be fainter than 26th magnitude (\\cite{big88}) - a result confirmed by HST observations (\\cite{big97}). In V Geminga is brighter at 25.4. This aspect of the spectrum has been explained by a proton cyclotron feature causing either preferential emission in V or absorption in B and I (\\cite{big96}) superimposed on a thermal continuum. However, re-analysis of the EUVE and ROSAT datasets highlight an error in this earlier work, indicating that the thermal continuum would not be expected to dominate in the optical regime, based on the observed flux (\\cite{hal96}). Such an apparent absorption feature has been previosuly observed in the Crab spectrum (\\cite{nas96}) although not confirmed by other observations (\\cite{kom96}). Recent spectral studies of Geminga (\\cite{mar98}) show a continuous power-law from 3700 to 8000 (\\AA) with no such features consequently indicating that a predominantly magnetospheric origin is preferred over a thermal one. It should be noted that these spectroscopic studies were at the limit of the observational capabilities of the Keck and with a low signal-to-noise ratio. Of crucial importance to the understanding of neutron star structure is the stellar radius. This can in principle be inferred once the distance and the black-body contribution has been measured (\\cite{wal97}). However determining the black-body component of an isolated neutron star is complicated by magnetospheric and possible atmospheric effects (\\cite{pav96}). As Geminga is very nearby it is a prime candidate for measuring the thermal component - crucial to this will be the removal of the magenetospheric component of its emission. This is possible by determining what contribution of the optical emission is pulsed and whether it 'follows' the hard (magnetospheric) or soft (presumed thermal) X-ray emission profile. The faintness of the optical counterpart has precluded time-resolved observations using conventional photometers. However by using 2-d photon counting detectors, the required astrometric analysis can be carried out off-line. Consequently photon arrival times can be measured from a reduced (seeing optimised) aperture diaphram. ", + "conclusions": "" + }, + "9802/astro-ph9802013_arXiv.txt": { + "abstract": "We present here the spectrum of the galaxy companion to the $z=4.7$ quasar BR1202-0725, in the optical range $6000 - 9000$\\AA, corresponding to $1050 - 1580$\\AA~ rest--frame. We detect a strong \\lya emission line at $z=4.702$, with an integrated flux of $2\\times10^{-16}$ergs cm$^{-2}$s$^{-1}$, and a UV continuum longward of the \\lya emission at a flux level of $\\simeq 3\\times10^{-19}$ergs cm$^{-2}$s$^{-1}$ \\AA$^{-1}$. We fail to detect any CIV$_{1550}$ emission with a $3 \\sigma$ upper limit of $3\\times10^{-17}$ergs cm$^{-2}$s$^{-1} $. We show that the ratio between \\lya and continuum intensity and the absence of a strong CIV emission imply that the UV continuum radiation is the result of an intense star--formation activity rather than of a reprocessing of the QSO flux. The total estimated SFR of this $z=4.7$ star--forming region is $\\sim 15-54$ M$_{\\odot}$ yr$^{-1}$, depending on the IMF and the metallicity. The present data suggest that the \\lya emission has a velocity and spatial structure, with possible velocity differences of $\\simeq 500$ \\kms on scales of few kpc. These velocity patterns may be a signature of collapsing or merging phenomena in the QSO and its environment, as expected from current models of galaxy formation at high $z$. ", + "introduction": " ", + "conclusions": "" + }, + "9802/astro-ph9802155_arXiv.txt": { + "abstract": "Old neutron stars (ONSs) which have radiated away their internal and rotational energy may still shine if accreting the interstellar medium. Rather stringent limits from the analysis of ROSAT surveys indicate that most optimistic predictions on ONSs observability are in excess of a factor as large as $\\sim 100$. Here we explore two possible evolutionary scenarios that may account for the paucity of ONSs. In the first it is assumed that the ONS population is not too fast ($V<100 \\ {\\rm km\\, s}^{-1}$) and that magnetic field decay guides the evolution. In the second, NSs move with high speed ($V>100$ km s$^{-1}$) and preserve their magnetic field at birth. We find that according to the former scenario most ONSs are now in the propeller phase, while in the latter nearly all ONSs are silent, dead pulsars. ", + "introduction": " ", + "conclusions": "" + }, + "9802/astro-ph9802296_arXiv.txt": { + "abstract": "New observations of SN 1980K made with the VLA at 20 and 6 cm from 1994 April through 1996 October show that the supernova (SN) has undergone a significant change in its radio emission evolution, dropping by a factor of $\\sim 2$ below the flux density $S\\propto t^{-0.73}$ power-law decline with time $t$ observed earlier. However, although $S$ at all observed frequencies has decreased significantly, its current spectral index of $\\alpha= -0.42\\pm0.15$ ($S \\propto \\nu^{+\\alpha}$) is consistent with the previous spectral index of $\\alpha=-0.60_{-0.07}^{+0.04}$. It is suggested that this decrease in emission may be due to the SN shock entering a new region of the circumstellar material which has a lower density than that expected for a constant speed ($w$), constant mass-loss rate ($\\mdot$) wind from the progenitor. If such an interpretation is correct, the difference in wind and shock speeds appears to indicate a significant evolution in the mass-loss history of the SN progenitor $\\sim 10^4$ years before explosion, with a change in circumstellar density ($\\propto \\mdot/w$) occurring over a time span of $\\lesssim 4 $ kyr. Such features could be explained in terms of a fast ``blue-loop'' evolutionary phase of a relatively massive pre-SN progenitor star. If so, we may, for the first time, provide a stringent constraint on the mass of the SN progenitor based solely on the SN's radio emission. ", + "introduction": "SN 1980K [RA(1950.0)$=20^{\\rm h}34^{\\rm m}26\\fs68 \\pm 0\\fs01;$ Dec(1950.0)$ = +59\\arcdeg55\\arcmin56\\farcs5 \\pm 0\\farcs2$] in NGC 6946 was discovered in 1980 October by Wild \\markcite{wild}(1980), and an initial, unsuccessful attempt was made on 1980 November 3 to detect it at 6 cm with the Very Large Array (VLA).\\footnote{The VLA is operated by the National Radio Astronomy Observatory (NRAO) of Associated Universities, Inc., under a cooperative agreement with the National Science Foundation.} However, only 35 days after optical maximum SN 1980K was detected at 6 cm on 1980 December 4 and has been regularly monitored at 20, 6, and occasionally 2 cm since then. Results from 1980 November through 1984 August are presented in Weiler et al.\\ \\markcite{w86}(1986), and results from 1984 November through 1990 December are presented in Weiler et al.\\ \\markcite{w92}(1992). By 1990 December the flux density had fallen to $S<0.3$ mJy at 20 and 6 cm, which made it difficult to observe with the VLA using typical, short ``snapshot'' observations, and regular monitoring was terminated. However, optical imaging and spectra taken after 1990 indicated that SN 1980K was still a detectable, albeit relatively constant optical luminosity source, particularly at H$\\alpha$ (Leibundgut \\markcite{l94}1994). Late-time optical and radio emission are predicted to be correlated (Chevalier \\& Fransson \\markcite{c94}1994), and in several examples are indeed observed to be related (see, e.g., Montes et al.\\ \\markcite{m97}1997, Van Dyk et al.\\ \\markcite{v98}1998 and references therein). Therefore, deeper and more sensitive VLA observations were made in 1994, 1995, and 1996 at both 20 and 6 cm, to attempt to obtain new measurements of SN 1980K below our previous sensitivity limits. Interestingly, these new data indicate that SN 1980K has, in fact, {\\em not} continued along its previously observed power-law decline, but has dropped sharply in radio flux density in the interval between 1990 and 1994. Additionally, recent measurements (Fesen \\markcite{f98}1998) indicate that SN~1980K has also faded in the optical over the last several years, possibly starting with a 20\\% decline by 1994 as reported by Fesen, Hurford, \\& Matonick (\\markcite{f95}1995). The new radio measurements are presented here, and they indicate a significant change in the evolution of the radio emission of SN 1980K from earlier, model-based expectations. ", + "conclusions": "At an age of $\\sim 10$ years, SN 1980K has apparently, and unexpectedly, entered a new phase of radio evolution, characterized by a large decrease in the flux density at both 6 and 20 cm below that expected from the model describing its previous power-law decline. The spectral index, however, is apparently unchanged and is consistent with its earlier value, indicating that the emission mechanism is likely the same. We interpret this radio decrease to imply that the circumstellar density has abruptly changed and the SN shock is now entering a CSM region $\\lesssim1.6$ times less dense than the RSG wind-established CSM previously encountered by the SN shock. This is consistent with models for an $\\sim 9 \\Msun$ star that undergoes a ``blue-loop'' evolutionary episode before entering its final RSG phase. It is interesting to note that changes in the radio emission, and, by implication, in the circumstellar density structure, have also been observed on time scales of a few years (a few $\\times 10^3$ years, in the time frame of the SN progenitor) for SN 1987A (Turtle et al.\\ \\markcite{t90}1990; Staveley-Smith et al.\\ \\markcite{s92}1992, \\markcite{s93}1993, \\markcite{s95}1995; Ball et al.\\ \\markcite{b95}1995; Gaensler et al.\\ \\markcite{g97}1997) and SN 1979C (Montes et al.\\ \\markcite{m98}1998). Although the radio emission from SN 1980K is now very faint ($\\lesssim 200\\mu$Jy), we intend to continue monitoring the SN with the VLA, to attempt to determine the reality, extent, and the radial profile for this proposed new CSM region. This may allow us to establish the properties and duration of this phase of pre-SN stellar evolution, and may, for the first time, provide a stringent constraint on the mass of the SN progenitor based solely on the SN's radio emission." + }, + "9802/astro-ph9802069_arXiv.txt": { + "abstract": "The high spatial resolution of the MECS experiment on board Beppo-SAX has encouraged a few scientists, including the author, to perform observations of galaxy clusters. Results from the analysis of the first few observed objects are encouraging. After having reviewed the Beppo-SAX observing program for galaxy clusters and referenced contributions to these proceedings by other authors on the same topic, I present results from the analysis of the Perseus cluster. ", + "introduction": "Although principally devoted to studies of compact galactic sources and active galactic nuclei, the Beppo-SAX satellite has devoted a small fraction of time to galaxy cluster observations. The primary motivation for performing observations of galaxy clusters with Beppo-SAX is the high spatial resolution of the MECS instruments. In these proceedings I briefly review the Beppo-SAX observational program for galaxy clusters. After having listed contributions on galaxy clusters to these proceedings by other authors I describe results from the analysis of the Perseus observation. ", + "conclusions": "" + }, + "9802/astro-ph9802188_arXiv.txt": { + "abstract": "RXTE has been operating for nearly 2 years and is planning the third. The spacecraft performance has been good and the three instruments are operating well. Observations have been made of the range of targets suitable for RXTE, including such different objects as accreting neutron stars and black holes, stellar flares, and supernova remnants. The goals of studying high time resolution and broad energy range and optimizing multi-wavelength participation are yielding important results. Oscillations found from low-mass X-ray binaries probably are signatures of the spin of the neutron stars and of the shortest orbital periods around the neutron stars. These are constraining neutron star parameters. Oscillation and spectral results from black hole candidates bring into the realm of possibility the possibilities of measuring the spins of the black holes and using X-ray data to test predictions of gravitation theory. Multi-wavelength observations are leading to identification of the locations of the X-ray emission regions and, in the case of the micro-quasars, to understanding of the mechanisms for jet formation. Recently faster observing response than originally planned has made possible some RXTE contributions to identification of gamma-ray bursts. ", + "introduction": "The RXTE was launched 1995 Dec 30. During the next week instruments were turned on while the spacecraft was checked out and calibrated. The Guest Observer (GO) program officially began Feb. 1, 1996, although a few time critical observations were done during the in-orbit check out period for the instruments. Some targets of opportunity, in particular, GRO J1744-28, were observed. The transient, bursting pulsar had been discovered by BATSE before RXTE's launch and interest in the unique object was high. Many Target of Opportunity (TOO) proposals had been accepted for the first 9 months of the GO program, but remarkably, this object was not envisioned and in fact there were no accepted proposals for a new, bright, pulsing transient. The User's Group had ruled that requests for observation of TOOs not covered by accepted proposals could be carried out if the data were made public. Thus we started early on a substantial public data base. Our planners make special effort to work new objects into the schedule and to find when sources can best be observed. In the two ensuing years 297 GO proposals were carried out, 71 of them (24 \\%) as TOOs. The amount of good observing time was $3.52 \\times 10^7$ s in 24 months, for a net efficiency of 56 \\%. RXTE was planned to have features judged necessary to understand the bright variable X-ray sky. It should have effective time resolution capable of studying the dynamical time scales of neutron stars and stellar black hole candidates, that is, better than tenths of milliseconds. For this, the area should be large. In fact, the area flown was cut back to be marginally effective for some goals. It should have energy coverage of at least 2-200 keV. Galactic compact sources in the plane tend to have column densities that cut off the observations at lower energies. Cyclotron lines, synchrotron radiation, and inverse-Compton spectra require sensitivity up to energies of $\\approx 100$ keV. Many spectra cut-off approaching $m_e c^2$, and it was impractical for us to achieve more than 200 keV within our constraints. Fast and flexible response would allow transient states to be studied and would facilitate multi-wavelength observations; these were also judged essential for this mission. \\begin{figure}[t] \\parbox{73mm}{\\epsfig{angle=0.,file=rome_mission_fig1.ps,width=73mm,height=73mm}} \\caption{ Drawing of RXTE in orbit. The five PCA counters are in the openings protected by solar shields. The 3 ASM Scanning Shadow Cameras rotate on the boom on the left. The HEXTE clusters are hidden under the thermal blankets to the right of the PCA. They look out the same direction as do the PCA collimators, with the same field of view. Two oppositely directed antennas allow scheduled communication to proceed regardless of changes in the observation schedule. } \\end{figure} In the launched configuration, the Proportional Counter Array (PCA) has an open area of 6250 cm$^2$ and the High Energy X-ray Timing Experiment (HEXTE) has 1600 cm$^2$ in phoswich scintillators. The Crab gives 13,400 counts s$^{-1}$ in the PCA, 284 counts s$^{-1}$ in HEXTE. The time resolution of the detector electronics and the data system of the PCA is $1 \\mu$s, for HEXTE, $8\\mu$s. There is a selection of data modes and steady telemetry rates of 50-256 kbps are regularly achieved, the lower value relatively steadily, the higher for periods of about a half hour, with the supplement of high rate dumps of the spacecraft solid state data recorder. These can be added by request with a days notice unless the Shuttle is flying. The achievements of high time resolution that I discuss below, show that this capability is successfully achieved. The responses and resolutions of the detectors are nominal, except that because of an electronic failure, for 1 of 8 HEXTE detectors the pulse heights are not available. HEXTE does have a larger than expected dead time due to large particle events, which reduces the sensitivity. Activation of the PCA during passage of the SAA has complicated calibration of the background. Calibrations of background and response will still be improved, but are now good to a couple of percent \\cite{Jahoda96,Jahoda98,Roth98}. Papers presented in this volume attest to the ability to use the PCA and HEXTE together to measure cyclotron line characteristics and the spectra of active galactic nuclei (AGN). Sources can be viewed with both the PCA and HEXTE when the bore-sights are farther than 30 degrees from the sun. This constraint is more relaxed than for most missions and has helped with the observations of transients and time critical behavior. Of course even this constraint can be frustrating. When GRO J1744-28 had a second outburst and for the second time BATSE saw that the first day's bursts were of a different character than later, with only 3 minutes separation time, we were unable to make observations. Many monitoring observations have followed sources for the duration of the mission. Brief looks (1000 s minimum) are possible for campaigns of various durations and density. Observations shorter than that, while sufficient to determine flux well, turned out to stress the satellite slewing and attitude determination resources. After the delivery of weekly observing plans, many re-plans are carried out to allow for TOOs and to change the observing plans on the basis of results. We had designed the system to the requirement that a new source could be observed within 7 hours of discovery. Within the regular planning process that is difficult, but such a rapid turn around is rarely optimal. However the Beppo-SAX discovery of the X-ray after-glows of Gamma-Ray Bursts has inspired development of much faster procedures \\cite{Marshall98}, which lead to observations within 2-3 hours. We can slew directly to a source (at 6 degrees per minute), if commanding is available and there are no conflicting stored commands. Figure 1 shows the satellite with the All Sky Monitor (ASM) on one end, comprising three one-dimensional imaging proportional counters with coded masks that rotate around the boom on which they are mounted and survey 70 \\% of the sky. Solutions for the flux of each cataloged source are computed for each 100 s dwell in the step-wise rotation per satellite orbit of 90 minutes. Residuals to the fit indicate uncataloged sources. The five detectors of the Proportional Counter Array operated flawlessly for about three months, at which time, the two detectors which operated at the coldest temperatures, within an orbit of each other suffered short-lived breakdown events. We have implemented the spacecraft's capacity to \"look\" at the data being produced and take prescribed actions to turn high voltage off in less than a minute if a breakdown event is beginning. This, and operating the spacecraft slightly differently to keep the detectors warmer, have so far allowed us to keep operating without accelerated detector degradation. During the first months the ASM detectors lost 4 anode wires out of 24, due to high voltage breakdown. The work-around in the ASM case was to burn the conducting carbon off the affected anodes and to guard against operating at very high rates. Each wire is calibrated \\cite{Rem97a,Levine98} and the results for a large number of sources can be found on the Web, on the MIT and the RXTE GOF pages: http://space.mit.edu/$\\sim$derekfox/xte/ASM.html http://heasarc.gsfc.nasa.gov/docs/xte. RXTEs targets include galactic compact stars, AGN, stars, supernova remnants, diffuse emission, and gamma-ray bursts. Galactic compact stars account for 80 $\\%$ of the targets: accreting high magnetic field binary pulsars, accreting low-mass binary pulsars, rotation-powered pulsars, black hole (BH) candidates, and white dwarfs. In this paper I will describe results that have been obtained for these sources. Other papers in this volume describe some of the results on AGN and diffuse sources. References must necessarily be incomplete and my choices of examples could well be biased by my own associations. ", + "conclusions": "RXTE has clearly been very successful for a variety of studies. The quantity and detail of the data make complicated systems like high magnetic field pulsars a challenge to analyze and complexity of the calibrations has daunted some observers initially. However, the results that have appeared are exciting and the good results from some observations of sources only a few mCrab in strength are inspiring. We look forward to many more anticipated results, to answering the deeper questions that observations have raised, and to observing more examples of the rare phenomena that we sometimes encounter that greatly improve our insight into high energy sources. Compared to the long run, the time these sources spend in some of the most interesting states is short, and there are few enough of them that the combination gives us relatively rare examples. We hope to use the hard won resources of RXTE and Beppo-SAX to be able to study as many of these as possible. I would like to acknowledge the many people who have contributed in many ways to planning, building, launching, operating, and using RXTE." + }, + "9802/astro-ph9802141_arXiv.txt": { + "abstract": "I argue that the rotation of white dwarfs is not a remnant of the angular momentum of their main sequence progenitors but a result of the mass loss process on the AGB. Weak magnetic fields, if present in stellar interiors, are likely to maintain approximately uniform rotation in stars, both on the main sequence and on the giant branches. The nearly uniform rotation of the core of the Sun is evidence for the existence of such fields. Exactly axisymmetric mass loss on the AGB from uniformly rotating stars would lead lead to white dwarfs with very long rotation periods ($>$ 10 yr). Small random non-axisymmetries ($\\sim 10^{-3}$) in the mass loss process, on the other hand, add sufficient angular momentum to explain the observed rotation periods around one day. The process illustrated with a computation of the probability distribution of the rotation periods under the combined influence of random forcing by weak nonaxisymmetries and angular momentum loss in the AGB superwind. Such asymmetries can in principle be observed by proper motion studies of the clumps in interferometric images of SiO maser emission. ", + "introduction": "White dwarfs are observed to rotate with typical periods of a day. The main sequence progenitors of these stars are also rotating, and it is generally assumed that the rotation of white dwarfs is a remnant of this main sequence rotation. Arguments involving conservation of angular momentum can be used to make this plausible (e.g.\\ Perinotto 1990, Pijpers 1993). A problem with this picture is, however, that progenitors of WD have gone through a giant stage in which at least the envelope rotated very slowly. Thus, it is necessary to assume that the cores of giants remain decoupled rotationally from their envelopes during the entire evolution from main sequence turnoff till the formation of the WD. Since little is known with certainty about the processes that might redistribute angular momentum inside stars, this assumption can not be easily rejected. There are, however, observational and theoretical reasons to doubt this picture. A strong observational argument is the internal rotation of the Sun. The most recent helioseismological measurements (Elsworth et al. 1995, Kosovichev et al. 1997, Corbard et al. 1997) show that the rotation below the convection zone is esentially uniform, with measured degrees of differential rotation well below the 30\\%. seen at the surface. The known hydrodynamic angular momentum transport processes, even with rather optimistic estimates of their efficiency, leave the Sun with a much too rapidly rotating interior (Spruit et al 1983). A new hydrodynamic mechanism recently studied in some detail is friction by internal gravity waves excited by the convection zone (Press 1981, Spruit 1987, Zahn 1990, Schatzman 1993). Realistic calculations of this process appear to be difficult, but estimates indicate that it can be more effective than the other hydrodynamic processes (Zahn et al. 1997). Magnetic fields, on the other hand, have long been known to be very efficient at transporting angular momentum. The torques exerted by magnetic fields become significant already at very low field strengths. For the Sun, for example, a field of less than 10G can provide sufficient torque to maintain the observed uniformity of rotation. A number of mechanisms can provide such weak fields, for example a fossil field (remnant of the star formation process) or a dynamo-like process operating on (a small remnant of) the differential rotation of the core. In this paper, I develop the consequences of assuming that the cores of giants do, in fact, corotate approximately with their envelopes. After discussing the observational evidence on WD rotation rates I develop theoretical arguments for the existence of effective coupling between the core and the envelope. This predicts very slowly rotating cores in the giant progenitor of a single WD. The rotation of single white dwarfs must then be explained by other processes. The same applies to neutron stars born in red giants. The observed pulsar rotation periods of the order of a second are much shorter (by a factor $10^3$ or so) than expected if they formed in approximately uniformly rotating giants, and with our assumption of strong coupling of the core another mechanism also has to be found to explain the rotation of pulsars. The processes differ somewhat for white dwarfs and neutron stars. The arguments for the neutron star case are developed in a companion paper (Spruit and Phinney, 1998). There, we show that the kicks with which neutron stars are born (as inferred from their transverse velocities) also impart angular momentum at amounts sufficient to explain the rotation of most pulsars. To explain the typical rotation rates periods of single white dwarfs (of which only about 30 have measured rotation rates), I show in Sect.\\ \\ref{nonax} that small asymmetries in the mass loss processes during the last phases of evolution on the AGB are sufficient to explain the observed rotation rates. These asymmetries act as a random forcing through which angular momentum accumulates in the envelope. A balance results between this random forcing and the loss of angular momentum by the wind. The evolution of the angular momentum as the mass loss proceeds turns out to be mathematically the same as that of the velocity of a particle experiencing Brownian motion in a gas, and can be described by a Fokker-Planck equation. Solutions of this equation (Sect.\\ \\ref{solu}) show that probability distribution of the angular momentum is close to a Maxwellian. The mean angular momentum decreases as the square root of the envelope mass remaining. Current observational evidence relating to the asymmetries needed in this picture is discussed in Sect.\\ \\ref{discu}. \\subsection{Rotation speeds of AGB cores} Starting with a rapidly rotating main sequence star, and assuming uniform rotation during the expansion to the giant stage, we can estimate the rotation periods to be expected for white dwarfs evolving from single stars. An early type main sequence (MS) star, rotating near its maximum speed (of the order 400 km/s), and expanding without angular momentum loss onto the asymptotic giant branch (AGB), has a rotation period $P_{\\rm G}=2\\pi/(GM/R_{\\rm G}^3)^{1/2} (R_{\\rm G}/R_{\\rm MS})^{1/2} \\sim 10$yr for $R_{\\rm G}\\sim 1$AU (except for a modest difference in gyration radius neglected here). Most early type MS stars rotate significantly slower, so that periods of the order 30-100 yr would be expected for the AGB descendants of early type stars. Some of the observed white dwarfs must have descended from solar type stars (F-G), which have periods of the order 30d at the end of their main sequence life. The AGB progenitors of these WD would rotate 100 times slower, with periods of the order of a thousand years. If the small amount of envelope mass is ignored which settles back onto the core during post-AGB evolution (more about this in Sect.\\ \\ref{axisy}), these rotation periods would also be inherited by the WDs formed. While there are a few magnetic white dwarfs with inferred rotation periods of at least a century, most WDs for which periods are known rotate much faster. We evidently need another mechanism to explain the rotation of typical single WDs. Before entering the discussion of possible mechanisms, I briefly review the observational evidence on WD rotation. ", + "conclusions": "\\label{discu} In the standard view, the rotating single WD derive from the rotating cores of giants, which somehow avoided spinning down in the slowly rotating convective envelope. I argue, instead, that rotating cores in giants are an unattractive idea, especially if these cores are magnetic. Unless the magnetic WD acquired their fields {\\it after} emerging from the envelope, the observed dipole moments are so large that a strong interaction with the slowly rotating convective envelope would be very hard to avoid. I recall the classical demonstration (e.g. in Mestel 1953, 1961) that rather weak magnetic fields (magnetic energy a small fraction of the rotational energy) can already transmit enough torques to maintain corotation between core and envelope. Such a weak field could be inherited from the star formation process. In order to prevent these torques from acting, any magnetic field in the core would have to be very weak or very accurately shielded from the convective envelope. In addition, a differentially rotating, initially nonmagnetic core is unstable to the growth of a small scale dynamo magnetic field, initiated by a magnetic shear instability (Balbus and Hawley 1992). The conditions for existence of this instability in stars were studied in detail already by Acheson (1978) who showed, in particular, that thermal diffusion allows it to operate under a much wider range of conditions than in the adiabatic case. The very weak differential rotation in the core of the Sun (e.g. Kosovichev et al.1997), for which no good explanation has been put forward except magnetic torques, is strong evidence for the operation of magnetic effects. While the arguments given here do not constitute a proof, I feel they are sufficiently compelling that approximately uniform rotation is a reasonable hypothesis, and is at least as plausible as the traditional assumption, which implies a core rotating $10^4$--$10^5$ times faster than the envelope for the entire duration of the RGB and AGB. I have explored the consequences of the assumption of approximately uniform rotation for AGBs stars in the process of shedding their envelopes. If this mass loss is strictly axisymmetric, the remaining core rotates very slowly (period more than 10 years). This is just the consequence of angular momentum conservation: the wind takes away almost the entire envelope, but the specific angular momentum it carries away is that of the stellar photosphere, which is larger than the average specific angular momentum of the envelope. On the other hand, only small nonaxisymmetries in the mass loss process suffice to give the star enough `kick' to explain the angular momentum of single white dwarfs. Such kicks could be associated with mass loss events at the pulsation period of the star or dust-formation episodes in the atmosphere. I have illustrated this with a calculation of the evolution of the probability distribution of the star's angular momentum under the combined action of many small nonaxisymmetric kicks and the angular momentum loss in the wind. The degree of asymmetry required is found to be of the order $10^{-3}$. Present theories for AGB mass loss are not detailed enough to calculate such asymmetries, but observational indications for asymmetries exist. Interferometric images of red supergiants ($\\alpha$~Sco, $\\alpha$~Ori and $\\alpha$~ Her: Tuthill, Haniff and Baldwin 1997), speckle reconstructions ($\\alpha$~Ori: Kluckers et al.\\ 1997) and HST imaging ($\\alpha$~Ori: Gilliland and Dupree, 1996) show pronounced `hot spots' on their surfaces. Assuming that such nonuniform photospheric conditions persist during the superwind phase, one would expect them to also affect the dust formation that is thought to be essential for the driving of the wind. The required asymmetry is obtained if a few (5 say) such spots are present, and the wind locally generated above these spots is slightly non-radial by a few tenths of a degree. That the mass flow is indeed asymmetric already close to the stellar photosphere is shown by speckle imaging (IRC 10216: Osterbart et al. 1997), and especially by mm-wave interferometric images of the SiO maser emission. These show a highly clumpy and time dependent structure (Diamond et al. 1994, Humphreys et al. 1996, Pijpers et al. 1994). This maser emission occurs at a few stellar radii, which is also the region where the backreaction of the wind on the star (`kick') takes place. Though the SiO maser emission is very sensitive to small changes in the local physical conditions, models of the emission (Lockett and Elitzur 1992, Bujarrabal 1994) should give estimates of the degree of inhomogeneity in the physical conditions in the wind. Measurement of deviations from radial flow in proper motion studies of the masing clumps in the wind should enable direct determination of the asymmetries relevant for the kick process described in this paper. An issue mentioned here only briefly is the origin of the 5 or so very slowly rotating ($P\\gapprox 100$yr) magnetic white dwarfs. A possible explanation is angular momentum loss in a radiation driven, but magnetized, wind during post-AGB evolution. This possibility will be further explored elsewhere. The coupling between core and envelope proposed here would also imply that the cores of pre-supernovae on the giant branch are so slowly rotating that very slowly ($P\\sim 1$hr) rotating neutron stars would result even if angular momentum were conserved during core collapse. While these would not show up as pulsars, one would have to argue that none of the observed pulsars were formed in red giants, which feels like an unattractive idea. It turns out, however, that the kicks neutron stars receive at birth and which give them their high observed space motion, are strong enough to impart a significant rotation as well. This idea is developed further in a separate paper (Spruit and Phinney, 1998)." + }, + "9802/astro-ph9802231_arXiv.txt": { + "abstract": " ", + "introduction": "Standard cosmology poses three basic conundrums: why the universe is homogeneous on large scales, why the cosmological density parameter $\\Omega_m$ is close to unity, and what seeds the large-scale structure in the universe. Inflation provides compelling solutions to all these problems (\\cite{Kolb&Turner}). The cosmic microwave background radiation (CMB) is observed to be isotropic to an accuracy of better than $10^{-4}$. This implies that the Universe was homogeneous to this accuracy over many horizon scales at last scattering. Within the standard cosmology there is no causal process which could have created the observed homogeneity. The basic idea of inflation is that the observable Universe grew from an initial patch small enough to fit inside the horizon. It is hypothesized that there were particle physics mechanisms which led to the rapid expansion of the early Universe required to produce the huge growth factor necessary to solve the cosmological problems above. In this paper, we assume that inflation produces primordial adiabatic Gaussian random-phase density fluctuations due to quantum fluctuations in the inflaton field. The properties of these fluctuations are completely specified by their power spectrum $A_S^2(k)$\\footnote{Certain two-field inflationary models predict isocurvature as well as adiabatic fluctuations (\\cite{Kof87}). Since the CMB and LSS predictions of isocurvature models differ significantly from adiabatic models and since these differences are not degenerate with parameter variation, we could, in principle, also fit for an isocurvature power spectrum. However, in this paper, we limit ourselves to considering adiabatic fluctuations.}. This paper will emphasize how well the shape of the primordial power spectrum can be measured from CMB anisotropy and redshift survey data. There currently exists a broad range of inflationary models (\\cite{Kolb96,Turner97}). Some of these models predict power spectra that are almost exactly scale-invariant (\\cite{Linde83}), or are described by a power law with spectral index less than one (\\cite{naturalinf,extendedinf}), while others predict power spectra with slowly varying spectral indices (\\cite{Wang94}), or with broken scale invariance (\\cite{Holman91ab,Adams97,Les97}). Indeed, there are physical reasons to believe that the primordial power spectrum has breaks in its power-law form. In effective (i.e., not complete) theories with two scalar fields, inflation may occur in two stages. The two stages of inflation can generate density perturbations with different power-law indices, with a step in amplitude of the primordial power spectrum on the scale of the transition between the two phases (\\cite{Holman91ab}). Fry \\& Wang (1992) found that such models can have a significant signature on small-scale CMB anisotropies. Recently, Adams, Ross, \\& Sarkar (1997) proposed a new multiple inflation model. Since attempts at an unified description of the strong, weak and electromagnetic interactions usually involve several stages of spontaneous symmetry breaking, they considered the effects of such symmetry breaking during an era of inflation in supergravity models. They found that there can be a succession of short bursts of inflation; the density perturbations produced during each burst is nearly scale-invariant but with differing amplitudes, and between each burst there is a brief period during which scale-invariance is badly broken. Given the range of possibilities for the primordial power spectrum, we would like to quantify how well it can be measured generically. Thus, we take the primordial power spectrum to be a free function in this paper. The upcoming data from the Microwave Anisotropy Probe (MAP; cf., Bennett \\etal\\ 1997; {\\tt http://map.gsfc.nasa.gov}) and the Sloan Digital Sky Survey (SDSS; cf., Gunn \\& Weinberg 1996) provide a unique opportunity for constraining inflationary models. It is useful to expand the temperature fluctuations in the CMB into spherical harmonics: $\\delta T/T (\\hat{\\bf r})= \\sum_{l,m} a_{T,lm} Y_{lm}(\\hat{\\bf r})$, where $\\hat{\\bf r}$ is the unit direction vector in the sky. MAP measures the angular power spectrum (\\cite{SeljakZ96}) \\be \\label{C_l def} C_{Tl} \\equiv \\langle |a_{T,lm}|^2 \\rangle= (4\\pi)^2 \\int \\frac{dk}{k}\\, A_S^2(k)\\, \\left| \\Delta_{Tl} (k, \\tau=\\tau_0) \\right|^2, \\ee where $A_S^2(k)$ is the power spectrum of the primordial density fluctuations (defined such that $A_S^2(k)=1$ for the scale-invariant Harrison-Zel'dovich spectrum), $\\Delta_{Tl} (k, \\tau=\\tau_0)$ is an integral over conformal time $\\tau$ of the sources which generate the CMB fluctuations, and $\\tau_0$ is the conformal time today. The power spectrum of mass fluctuations in the linear regime today is \\be P(k)=P_0 \\,k A_S^2(k)\\, T^2(k), \\ee where $P_0$ is a normalization constant, and $T(k)$ is the transfer function, which depends on physics at matter-radiation equality and decoupling. Note that $P_0$ has been absorbed into the definition of $\\Delta_{Tl} (k, \\tau=\\tau_0)$ in Eq.~(\\ref{C_l def}). The galaxy redshift survey from the SDSS (cf., Strauss 1997 for a description) will allow a determination of $P_G(k)$, the galaxy power spectrum in redshift space. The quantity $P_G(k)$ differs from the mass power spectrum $P(k)$ due to two effects. First, the galaxy distribution may be biased with respect to the mass distribution. On large scales, models (Weinberg 1995; Kauffmann, Nusser, \\& Steinmetz 1997; Scherrer \\& Weinberg 1997) indicate that the mass and galaxy density fields are directly proportional to one another. The proportionality constant is referred to as the galaxy bias parameter, $b$. Second, peculiar velocities cause the density contrast of galaxies in redshift space to appear appreciably stronger than in real space. In linear theory, the power spectrum is boosted by a factor $1 + {2 \\over 3} \\beta + {1 \\over 5} \\beta^2$, where $\\beta \\equiv \\Omega_m^{0.6}/b$ ($\\Omega_m$ is the matter density in units of the critical density of the Universe $\\rho_c \\equiv 3H_0^2/(8\\pi G)$, with $H_0$ denoting the Hubble constant). The net result (Kaiser 1987; cf., Hamilton 1998 for a review) is that on linear scales, the galaxy power spectrum is given by: \\begin{equation} \\label{eq:PG} P_G(k) = b^2 \\left(1 + {2 \\over 3} \\beta + {1 \\over 5} \\beta^2\\right) P(k) \\equiv b_{\\it eff}^2 P(k). \\end{equation} The observables $C_{Tl}$ and $P_G(k)$ depend on the cosmological parameters $H_0$, $\\Omega_b$ (baryon density in units of $\\rho_c$), $\\Omega_m$, $\\Omega_{\\Lambda}$ (density contribution from the Cosmological Constant in units of $\\rho_c$), and $\\tau_{ri}$ (reionization optical depth) only through $\\Delta_{Tl} (k, \\tau=\\tau_0)$, $b_{\\it eff}$, and $T(k)$. The MAP data alone will allow rather tight constraints on the baryon/photon ratio (which determines $\\Omega_b h^2$), the matter/photon ratio (which determines $\\Omega_m h^2$), and the geometry of the universe, and test the basic inflationary scenario (Spergel 1994; Knox 1995; Jungman \\etal\\ 1996; Bond, Efstathiou \\& Tegmark 1997; Zaldarriaga, Spergel \\& Seljak 1997; Dodelson, Kinney \\& Kolb 1997; Lidsey \\etal\\ 1997; Copeland, Grivell \\& Liddle 1997; Efstathiou \\& Bond 1998; Eisenstein, Hu \\& Tegmark 1998). However, there is a near-degeneracy in several sets of parameters (Bond \\etal\\ 1994; Bond, Efstathiou \\& Tegmark 1997; Zaldarriaga, Spergel \\& Seljak 1997; \\cite{HD97}) including the overall amplitude of the power spectrum at $k \\sim 0.1\\,h\\rm Mpc^{-1}$ and the optical depth. Since $\\Delta_{Tl} (k, \\tau=\\tau_0)$ and $T(k)$ depend on the cosmological parameters differently, by combining the MAP and SDSS data, we can break these degeneracies and determine cosmological parameters to impressive accuracies. Several other recent papers have also considered how well cosmological parameters can be constrained with the combination of future CMB and galaxy survey data: Tegmark (1997) has studied a somewhat different and smaller set of the cosmological parameters than we consider here; Hu, Eisenstein, \\& Tegmark (1997) and Eisenstein, Hu, \\& Tegmark (1998) have tested the sensitivity of the data to the neutrino mass, which we do not consider in this paper. See also Lineweaver \\& Barbosa (1998) and Webster \\etal\\ (1998), who fit for cosmological parameters using the best available CMB and redshift data available. Both $C_{Tl}$ and $P_G(k)$ are proportional to the amplitude of the power spectrum of the primordial density fluctuations $A_S^2(k)$. Because of the finite resolution of the MAP satellite, the errors on the $C_{Tl}$ increase rapidly with wavenumber for $l \\ga 800$, while the errors on $P_G(k)$ from the SDSS decrease with wavenumber; MAP and SDSS thus quite naturally complement each other in the determination of $A_S^2(k)$. Assuming that inflation occurred, all the obtainable information about the inflationary model is contained in $A_S^2(k)$. Some inflationary models predict a significant gravity wave contribution to the microwave background fluctuation spectrum at small $l$. However, even in these models, gravity waves do not make a significant contribution to the microwave background fluctuation spectrum at $l \\ga 100$ (\\cite{Crit93,Dode94,Allen95,Wang96,Turner96}). Since most of the information on the cosmological parameters comes from relatively small angular scales ($l \\ga 100$), the inclusion of tensor perturbations would not change our results qualitatively. We parameterize $A_S^2(k)$ by a series of steps equally spaced in $\\log k$ bins, each step with a constant amplitude in its bin. By taking these amplitudes to be independent parameters, we probe features in the primordial power spectrum on all scales. The measurement of the primordial power spectrum independent of specific inflationary models from the combined MAP and SDSS data should shed light on our understanding of the early Universe. In Section 2, we present the statistical technique we use in estimating errors of parameters. Before considering a general primordial power spectrum, we determine the accuracy with which cosmological parameters will be measured with a pure power law $A_S^2(k)$ in Section 3. In Section 4, we relax this constraint, and discuss the determination of the primordial power spectrum, using the stepwise parameterization described above. Section 5 contains a summary of the paper. ", + "conclusions": "Cosmic microwave background experiments measure fluctuations in the curvature of space at the surface of last scattering. If the primordial fluctuations are adiabatic, then MAP will be a sensitive probe of the geometry of the universe, the matter/photon ratio ($\\Omega_m h^2$), the baryon/photon ratio ($\\Omega_b h^2$) and the primordial power spectrum for $k < 0.2\\,h\\rm Mpc^{-1}$. Observations of large scale structure measure the large scale distribution of galaxies. If the primordial fluctuations are adiabatic, then SDSS will be a sensitive probe of the primordial power spectrum for $k > 0.02\\,h\\rm Mpc^{-1}$ and $\\Omega_m h$, which determines the horizon size at matter-radiation equality and thus the shape of the processed power spectrum. SDSS will also be a powerful probe of the nature of the dark matter (Hu \\etal\\ 1997). Because of the large number of galaxies with detailed photometric and spectroscopic data in the SDSS, we can determine the power spectrum and bias factor for several independent samples of galaxies, defined by morphology, color, or spectral characteristics. To the extent to which these independent samples yield consistent power spectra, we can determine the range in $k$ over which the linear bias model is valid. The combination of the SDSS and the MAP data will be a powerful test of our basic cosmological models. Both experiments will accurately determine the amplitude of density fluctuations at $0.2 > k > 0.02\\, h\\rm Mpc^{-1}$. Will there be a set of cosmological parameters and primordial power spectrum that is consistent with both sets of observations? If so, then we will have tested the gravitational structure formation paradigm, our interpretation of the primordial fluctuations as adiabatic, and the linear biasing paradigm. If not, this conflict will likely lead us to a deeper understanding of the origin of structure in the universe. In particular, from the combination of the MAP and SDSS data, we can obtain a measurement of the primordial power spectrum without reference to any specific inflationary model, assuming that the primordial fluctuations are adiabatic. If the values of the cosmological parameters that best fit the data are close to what we expect, but the primordial power spectrum differs significantly from the predictions of the current inflationary models, then it will be an indication of new physics in the early universe, and it will provide a solid starting point for building new inflationary models. If there is a theory consistent with both data sets, then the combination of the two observations will break the degeneracies inherent in each individual observation. Cosmic microwave background observations by themselves measure a combination of the amplitude of the primordial power spectrum and the optical depth. Large scale structure observations measure the product of the current matter power spectrum and the bias parameter. In linear theory, the offset between the power spectrum as determined by the two sets of observations determines a combination of the bias parameter and the optical depth of the universe. Because of non-linear effects on the spectrum, this offset is scale-dependent, this will enable us to independently measure the bias parameter and the optical depth of the universe. The combination of the two experiments will also improve our ability to determine other cosmological parameters, in particular, the matter density and the cosmological constant. Measurements of the cosmic microwave background polarization can provide an independent measure of the optical depth of the universe $\\tau_{ri}$ (Zaldarriaga, Spergel \\& Seljak 1997). However, emission from both polarized galactic dust and synchrotron emission (Keating \\etal\\ 1998) may swamp the primordial polarization fluctuations at large angular scales. These foregrounds may significantly limit our ability to extract useful cosmological information from polarization measurements. Hence it is important that we can use the SDSS data as an alternative and independent aid to MAP in the determination of $\\tau_{ri}$. The Sloan Digital Sky Survey, and other large scale structure surveys, will make other independent measurements that will probe cosmological parameters and the primordial spectrum: redshift distortions, cluster properties, small scale velocity fields, and the evolution of structure. These will provide additional tests of the basic model and will further enhance our ability to measure cosmological parameters. With this in mind, there are a number of improvements that could be done in this analysis. The redshift distortions could be measured directly, and the effects could be included directly into the Fisher matrix analysis (cf., de Laix \\& Starkman 1998 and Hatton \\& Cole 1998 for a discussion of how well redshift-space distortions can be measured from the SDSS data). Similarly, we could parameterize the effects of galaxy and clustering evolution, and include these as parameters in the analysis. More challenging will be a proper accounting of non-linear effects on small scales, in the power spectrum, the redshift space distortions, and the bias (here assumed linear and independent of scale). The analyses in Tables~2-4 show that going to $k_{max} = 0.5\\,h\\rm Mpc^{-1}$ gives us particularly strong constraints on cosmological parameters, but we have taken the non-linear effects into account in only a relatively crude way in this paper. It remains to be seen the extent to which the uncertainties due to these inaccuracies strongly affect the error bars we derive on the primordial power spectrum. The next few years will be a very exciting time in cosmology." + }, + "9802/astro-ph9802095_arXiv.txt": { + "abstract": "We are using the 2dF spectrograph on the Anglo-Australian Telescope to obtain spectra for a complete sample of all 14000 objects with $16.5$ 1.0 \\msun) stars, but there may still be too many massive white dwarfs to be explained by the predictions of single star evolution, with a conventional initial mass--final mass relation, initial mass function, and star formation history. Recent theoretical work (\\cite{yungetal}, hereafter YLTS; \\cite{ibtutyung}, hereafter ITY) reveals that consideration of the \"visibility function\", a consequence of WD cooling along evolutionary sequences, is an important factor in calculations of the expected DD number distribution. Systems born with orbital periods $P$~\\gta 10$^{\\rm h}$ do not evolve to shorter periods by emission of gravitational radiation on interesting (i.e., short) timescales. Below 10$^{\\rm h}$, an approximate equilibrium is established between systems born into short-period orbits and those that vanish either by merging or by cooling in some \\tento{8} years beyond the detection limits of surveys with relatively bright limiting magnitudes, such as our own. Only a small fraction of these pairs will actually merge in the \\tento{8} year period of visibility, and only \\twid 1/20 of the pairs with $M_{\\rm tot}$ \\gta \\mch\\ are visible (cf. ITY, Figure 2). These results indicate that previous DD searches had sample sizes too small to reveal the massive, short-period systems of interest. On the other hand, it is clear that the close DDs that {\\em were\\/} discovered (Saffer et al. 1988; BGRD; \\cite{tmarsh1}; \\cite{tmarsh2}) fall within the peak of the expected DD number distribution (cf. YLTS, ITY). YLTS suggested that further searches with increased sample sizes are likely to reveal the existence of the short-period, massive DD population. In this paper, we report observations of an enlarged sample in order to place meaningful constraints both on DDs as \\snia\\ progenitors and, more generally, on the theory of close binary star evolution. ", + "conclusions": "We have performed a radial velocity survey of a total of 153 field WDs and sdBs, most of which were not previously known to be binary. In the combined sample, we have discovered 18 new DD candidates with WD or sdB primaries, and 1 new confirmed WD+MS pair. Among the 7 sdB variables, we have obtained orbital solutions for the short-period sdB+WD pairs Feige 36 ($P$ = 8\\fh 5) and Ton 245 ($P$ = 2\\fd 5), and Moran et al. (1998, in preparation) have found orbital periods for 0101+039 ($P$ = 16\\fh 1 or 13\\fh 7), 1432+159 ($P$ = 5\\fh 39), and 2345+318 ($P$ = 5\\fh 78). Our conclusions are: \\begin{enumerate} \\item The population of close, short-period DD systems predicted by the theory of close binary star evolution does in fact exist in significant numbers. The raw observed fractions of confirmed short-period DDs and WD+MS pairs among all WDs are in satisfactory agreement with the predictions of theory for values of the common envelope efficiency parameter \\ace\\ near 1.0. However, applying a correction for selection effects implies an excess of existing DDs over the predicted number. A resolution of this apparent discrepancy will have to await a more complete understanding of the CE process. \\item The observed DD orbital period and primary mass distributions do fall very nearly in the peak of the predicted distributions for \\ace\\ = 1.0, although possibly shifted slightly. We note that the discrepancy of the predicted and observed period distributions could be made to vanish by a small decrease in the value of \\ace, a free parameter in the theory (``the parameter of ignorance'', I. Iben). The discrepancy with the predicted mass distribution is less amenable to fine-tuning of the theory, and we suggest that some previous primary WD mass determinations may have been systematically biased to higher masses due to dilution of the optical spectra by the presence of unseen, higher-gravity companions. Over-representation of sdB primaries in the combined sample is responsible for the apparent excess of stars with masses near 0.5 \\msun. A refined analysis of the effect of the mass-dependence of WD cooling sequences upon the $\\log P$--$\\log M$ number distribution is also required. \\item The low-mass systems observed by Marsh et al. (1995) and Marsh (1995) tend to have short orbital periods, in agreement with theoretical predictions. However, their total system masses likely are too small to make them viable \\snia\\ progenitors. Furthermore, the primary stars must have He cores, while CO white dwarfs are thought to be the best candidates to undergo the central ignition and detonation of carbon required to produce a \\snia. \\\\ \\item The one known super-\\mch\\ DD system, LB~11146, has an unknown orbital period, although its projected orbital separation is \\lta 1 a.u. The large remnant masses suggest that the progenitor stars also were massive, had large radii in the AGB phase, and that the system might have gone through at least one CE phase leading to a short orbital period. However, no ``loaded gun'' has yet been found, i.e., a super-\\mch\\ system with $P$ \\lta 10$^{\\rm h}$. On the other hand, the sample of confirmed short-period DDs is still smaller than the number predicted to contain a massive, short-period system. When confirmed, the 18 new DD candidates discovered in our survey will increase the number of known short-period DDs among all WDs to a fraction approaching that which theory predicts will include a pre-\\snia\\ system. \\\\ \\end{enumerate} We shall continue with follow-up observations to solve for the orbital parameters of the new WD+WD and sdB+WD candidates found in our survey, and with the survey work itself to further expand the sample. Additional work on the DAO WDs is needed to elucidate the connection between close binary evolution and the very low masses of half of DAOs. Infrared observations, especially for the higher-luminosity sdB stars, will be required to distinguish main sequence companions having masses as low as \\twid 0.1 \\msun\\ from compact companions. Our survey has revealed a significant new population of close DD systems. Even if no convincing potential \\snia\\ progenitor is found, the distributions of DD orbital periods and primary masses (and secondary masses and mass ratios, where observeable) constitute powerful probes of the CE process, and a meaningful test of the predictions of the theory of close binary evolution now seems within reach. \\bigskip Acknowledgements: This work has been supported by NASA Grants NAGW-2678, G005.44000, and by the Director's Discretionary Research Fund at the Space Telescope Science Institute. LRY acknowledges support through RFBR grant 960216351, and is grateful for the hospitality of Meudon Observatory. RAS acknowledges fruitful conversations with Tom Marsh, Jim Liebert, Jay Holberg, and Gary Schmidt. LRY acknowledges day-to-day discussions of stellar evolution with Alexander Tutukov. The authors are grateful for the long-term committment of the KPNO TAC and awards of valuable Kitt Peak 2.1-m time, crucial to the success of our large-scale survey. We thank the anonymous referee for a critical review of the manuscript. \\clearpage" + }, + "9802/astro-ph9802160_arXiv.txt": { + "abstract": "We report on the results of optical follow-up observations of the counterpart of GRB 970508, starting 7 hours after the event. Multi-color U, B, V, R$_{\\rm c}$ and I$_{\\rm c}$ band observations were obtained during the first three consecutive nights. The counterpart was monitored regularly in R$_{\\rm c}$ until $\\sim$ 4 months after the burst. The light curve after the maximum follows a decline that can be fitted with a power law with exponent $\\alpha$ = --1.141 $\\pm$ 0.014. Deviations from a smooth power law decay are moderate (r.m.s. = 0.15 magnitude). We find no flattening of the light curve at late times. The optical afterglow fluence is a significant fraction, $\\sim$ 5\\%, of the GRB fluence. The optical energy distribution can be well represented by a power law, the slope of which changed at the time of the maximum (the spectrum became redder). ", + "introduction": "\\label{sec:intro} With the Wide Field Cameras (WFCs; Jager et al. 1995) onboard BeppoSAX (Piro et al. 1995) it became possible for the first time to rapidly determine accurate (3\\arcmin\\ radius) GRB positions (Costa et al. 1997a,b; In 't Zand et al. 1997; Heise et al. 1997a); such positions, as it turned out, are the necessary ingredient for succesful counterpart searches (see also Takeshima et al. 1997 and Smith et al. 1997). The detection of a fading optical counterpart to GRB 970228 provided the first arcsecond localization of a GRB (Groot et al. 1997a; Van Paradijs et al. 1997; Galama et al. 1997a,b). Its coincidence with a faint extended object (Groot et al. 1997b; Metzger et al. 1997a; Sahu et al. 1997a), which is probably the host galaxy, suggests a cosmological distance for the GRB. A second optical transient (OT) was found for GRB 970508 (Bond 1997). The measurement of a lower bound to the redshift ($z \\geq$ 0.835; Metzger et al. 1997b) of this optical counterpart has unambiguously settled the cosmological distance scale to GRB sources. On May 8.904 UT the Gamma-Ray Burst Monitor (GRBM) on BeppoSAX recorded a moderately bright GRB (Costa et al. 1997c), which was also recorded with its Wide Field Cameras. The fluences, recorded with the GRBM (40-700 keV) and WFC (2-26 keV), were (1.8 $\\pm$ 0.3) $\\times 10^{-6}$ erg\\,cm$^{-2}$ and (0.7 $\\pm$ 0.1) $\\times 10^{-6}$ erg\\,cm$^{-2}$, respectively (Piro et al. 1997a). The burst was also recorded (Kouveliotou et al. 1997) with BATSE (Fishman et al. 1989) with a total 20-1000 keV fluence of $(3.1 \\pm 0.2) \\times 10^{-6}$ erg\\,cm$^{-2}$, and a peak flux density (50-300 keV) of ($1.66 \\pm 0.06) \\times 10^{-7}$ erg\\,cm$^{-2}$\\,s$^{-1}$. From optical observations of the WFC error box (Heise et al. 1997b), made on May 9 and 10, Bond (1997) found a variable object at RA = $06^{\\rm h}53^{\\rm m}49\\fs2$, Dec = +79\\degr16\\arcmin19\\arcsec (J2000), which showed an increase by $\\sim$1 mag in the V band. BeppoSAX Narrow Field Instrument observations revealed an X-ray transient (Piro et al. 1997b) whose position is consistent with that of the optical variable. Extended emission is not associated with this optical counterpart (Fruchter et al. 1997, Sahu et al. 1997b, Pian et al. 1997), and Natarajan et al. (1997) showed that either GRB 970508 originated from an intrinsically very faint dwarf galaxy or that it occured at a large distance from a host galaxy ($> 25 h_{70}^{-1}$ kpc). We here report on the results of optical photometry of GRB 970508, made between 0.3 and 110 days after the burst. ", + "conclusions": "} The differential R$_{\\rm c}$ light curve shows that deviations from a smooth power law decay are moderate (r.m.s. = 0.15 mag; lower panel in Fig. \\ref{fig:lightcurve}). The magnitude $R_{\\rm c} = 26.09 \\pm 0.36$ for a possible underlying galaxy is consistent with the limits derived by Sahu et al. (1997b) and Pian et al. (1997), i.e., $R_{\\rm c} > 25.5$. We note, however, that the introduction of a constant to the power law fit has not improved the $\\chi^{2}$ substantially; we conclude that the evidence for a flattening of the light curve at late times is not yet convincing. The optical fluence $S_{\\infty}$ (3000\\AA-10000\\AA, $\\beta = -1$; Sect. \\ref{sec:spec}), is about 5\\% of the BATSE 20-1000 keV and 20\\% of the WFC 2-26 keV GRB fluences. A similar comparison with the X-ray afterglow (Piro et al. 1997a; $S_{X}$ = 7.3$\\times 10^{-7}$ erg\\,cm$^{-2}$ 2-10 keV, from an extrapolation of the X-ray light curve for $t$ $>$ 27 s) shows that the optical afterglow fluence is $\\sim$ 20 \\% of the X-ray afterglow fluence. The optical spectrum of the OT becomes redder before reaching maximum light (Fig. \\ref{fig:slopes}); afterwards the spectral slope remains constant. We distinguish three phases of the light curve: phase I, the rising phase ($\\alpha = 3.0\\pm0.4,\\beta=-0.33\\pm0.17)$, phase II at maximum light ($\\alpha \\sim 0,\\beta=-0.9\\pm0.10)$, and phase III, the decaying phase ($\\alpha = -1.141 \\pm 0.014,\\beta=-1.11\\pm0.06)$. Wijers et al. (1997) give a relation between the spectral slope $\\beta$ and the temporal slope $\\alpha$ for the simplest fireball remnant model (a forward blast wave with single Lorentz factor $\\Gamma$; their Eq. 2) and for a `beamed' fireball, i.e. $\\Gamma$ and the energy per unit solid angle E are functions of angle (their Eq. 4). For phase III we find very good agreement with $\\alpha= -1.141 \\pm 0.014$ implying $\\beta = -1.094 \\pm 0.009$ for the `beamed' case (marginal agreement with $\\beta = -0.761 \\pm 0.009$ for the simple case). Phases I and II are not in agreement with the simple and beamed fireball models, that predict less steep or even inverted spectral slopes $\\beta$. However, these early phases might be explained by a continuous distribution of Lorentz factors (Rees \\& M\\'esz\\'aros 1997). Lower $\\Gamma$ (redder) material takes longer to sweep up a significant amount of external matter and hence its emission is delayed with respect to the higher $\\Gamma$ (bluer) material. The lower $\\Gamma$ material catches up gradually with the decelerated higher $\\Gamma$ material. Hence the rise of the light curve and evolution from blue to red can be explained with building up the emission from material with lower $\\Gamma$, while the power law decay (phase III) sets in when most of the material has been mixed to a single $\\Gamma$." + }, + "9802/astro-ph9802210_arXiv.txt": { + "abstract": "Quantitative measures of central light concentration and star formation activity are derived from R and H$\\alpha$ surface photometry of 84 bright S0-Scd Virgo Cluster and isolated spiral galaxies. For isolated spirals, there is a good correlation between these two parameters and assigned Hubble types. In the Virgo Cluster, the correlation between central light concentration and star formation activity is significantly weaker. Virgo Cluster spirals have systematically reduced global star formation with respect to isolated spirals, with severe reduction in the outer disk, but normal or enhanced activity in the inner disk. Assigned Hubble types are thus inadequate to describe the range in morphologies of bright Virgo Cluster spirals. In particular, spirals with reduced global star formation activity are often assigned misleading early-type classifications, irrespective of their central light concentrations. 45$\\pm$25\\% of the galaxies classified as Sa in the Virgo Cluster sample have central light concentrations more characteristic of isolated Sb-Sc galaxies. The misleading classification of low concentration galaxies with low star formation rates as early-type spirals may account for part of the excess of `early-type' spiral galaxies in clusters. Thus the morphology-density relationship is not all due to a systematic increase in the bulge-to-disk ratio with environmental density. ", + "introduction": "Discerning environmental effects on the evolution of galaxies depends on objective comparisons between cluster and field galaxies. Traditionally, comparisons are drawn between galaxies that are assigned to the same Hubble class (e.g., Kenney 1990, Oemler 1992). However, the Hubble classification system was based mainly on nearby field galaxies, and may not be adequate to describe environmentally altered galaxies in dense environments. Attempts to force cluster galaxies into Hubble type morphological bins may lead to biased conclusions about the physical differences between cluster and field galaxies, particularly those of the same assigned Hubble type (see also van den Bergh 1997). Since the Hubble classification is the framework for the well-known morphology-density relationship (Oemler 1974; Dressler 1980), it is especially important to study whether misleading classifications contribute significantly to the increase in the fraction of S0's and early-type spirals (Giovanelli, Haynes, \\& Chincarini 1986) in clusters. As defined by Hubble (1936), classification for spiral galaxies depends on three different criteria: the relative brightness of the `nuclear' and spiral components (e.g., the bulge-to-disk-ratio, the degree of central light concentration), the degree to which HII regions are resolved in the spiral arms, and the openness of the spiral arms. There is a correlation in the mean between the three criteria for many nearby field spiral galaxies, although the scatter is large (Kennicutt 1981; Bothun 1982). Van den Bergh (1976) suggested that application of these criteria in a cluster environment leads to misleading galaxy classifications, particularly since the second criterion is strongly dependent on the star formation rate, which may be systematically affected in clusters (e.g., Kennicutt 1983, Moss \\& Whittle 1993, Koopmann \\& Kenney 1998b). Van den Bergh based his DDO classifications solely on apparent bulge-to-disk ratios (much as the Yerkes classification (Morgan 1958) is based primarily on a visual light concentration parameter) and introduced the term `anemic' to refer to galaxies with weak star formation. Bothun (1982) stressed the use of \\it quantitative \\rm parameters, such as the bulge-to-disk ratio, to better trace the underlying stellar mass distribution of the galaxy and avoid dependence on visually striking star formation characteristics, as well as resolution effects. Recent studies (Dressler et al. 1994; van den Bergh et al. 1996) have affirmed that a relatively smaller fraction of distant galaxies fits the Hubble classification, and several groups (e.g., Abraham et al. 1994; Fukugita et al. 1995; Hashimoto et al. 1998) have adopted quantitative parameters, such as central light concentration, to trace morphological type. However, no quantitative test of the Hubble classification has yet been made for nearby cluster galaxies. In this paper, we explore whether the Hubble types of nearby isolated and cluster galaxies correlate with quantitative measures of the central light concentration and star formation activity. This work is part of a series on the comparative star formation properties of Virgo Cluster and isolated galaxies. A full discussion of star formation properties in the two samples is given in Koopmann \\& Kenney (1998b). ", + "conclusions": "We have quantitatively shown that the weakness of star formation in some small to intermediate concentration Virgo Cluster galaxies causes them to be assigned early spiral or uncertain Hubble types, as suggested qualitatively by van den Bergh (1976). Cluster galaxies, particularly those with weak star formation, are not necessarily similar in central concentration to isolated galaxies with similar Hubble classifications. Data from the literature suggest that there is evidence for misleading classifications in other clusters, including Coma (Caldwell et al. 1996; Bothun, Schommer, \\& Sullivan 1983) and Abell 957 (Abraham et al. 1994). To what extent do misleading Hubble classifications contribute to the morphology-density relationship? While several studies have concluded that cluster S0 populations cannot be produced by stripped spirals (Dressler 1980; Boroson, Strom, \\& Strom 1983), others suggest they can (Solanes, Salvador-Sol\\'e, \\& Sanrom\\'a 1989; Eder 1990, Bothun \\& Gregg 1990). Our work shows that there is a significant population of low concentration spirals with reduced star formation in the Virgo Cluster. Thus the excess of `early-type' spiral galaxies in nearby clusters is partially due to misleading classifications of low concentration systems with reduced star formation rates, and not all due to a systematic increase in bulge-to-disk ratio with environmental density. These results are especially interesting in light of the work of Dressler et al. (1997), who suggest that as many as half the S0 galaxies in nearby clusters have evolved from galaxies which were spirals at z$\\sim$0.5. We propose that most of the low concentration `early-type' Virgo spirals with low star formation rates were blue, star-forming, `late-type' spirals a few Gyr ago, and that they have recently evolved due to cluster environmental processes. What environmental processes have contributed to the evolution of low concentration galaxies with low star formation rates? ICM-ISM stripping is likely to play an important role in the removal of gas and star formation from the outer disk; the Virgo Cluster spiral NGC 4522 is an example of ongoing stripping (Kenney \\& Koopmann 1998). Tidal interactions between galaxies, which cause the redistribution of gas inwards and/or outward, also are important in the Virgo Cluster. At least one of the low concentration Virgo Cluster galaxies with reduced outer disk star formation, the St galaxy NGC 4424, is a merger remnant, (Kenney et al. 1996) and others show evidence of disturbed outer stellar disks (Koopmann \\& Kenney 1998b). Whether one type of environmental process is dominant is not yet clear, and at least some galaxies may experience multiple environmental interactions. Funding for this work was provided by NSF grant AST-9322779. We are grateful to Y. Hashimoto, G.D. Bothun, S. Jogee, R.B. Larson, and V.C. Rubin for helpful comments." + }, + "9802/astro-ph9802032_arXiv.txt": { + "abstract": "Old neutron stars (ONSs) which have radiated away their internal and rotational energy may still shine if accreting the interstellar medium. Despite their large number, only two promising candidates have been detected so far and rather stringent limits on their observability follow from the analysis of ROSAT surveys. This contrasts with optimistic theoretical estimates that predicted a large number of sources in ROSAT fields. We have reconsidered the issue of ONSs observability, accounting for the spin and magnetic field evolution over the neutron star lifetime. In the framework of a spin--induced field decay model, we show that the total number of ONSs which are, at present, in the accretion stage is reduced by a factor $\\sim 5$ over previous figures if the characteristic timescale for crustal current dissipation is $\\sim 10^8-10^9$ yr. This brings theoretical predictions much closer to observational limits. Most ONSs should be at present in the propeller phase and, if subject to episodic flaring, they could be observable. ", + "introduction": "Old neutron stars (ONSs), i.e. neutron stars which have evolved beyond the pulsar phase, are expected to be quite a large Galactic population counting as many as $N\\sim 10^9$ objects. Ostriker, Rees, \\& Silk (1970)\\markcite{ors70} were the first to suggest that accretion of the interstellar medium (ISM) may produce enough luminosity to make the closest stars observable and that ONSs may contribute to the soft X--ray background. More recently, Treves, \\& Colpi (1991)\\markcite{tc91} and Blaes, \\& Madau (1993)\\markcite{bm93} have shown that several thousands ONSs should be present in the ROSAT All Sky Survey. Rather optimistic predictions for the number of detectable sources were also presented by Madau, \\& Blaes (1994)\\markcite{mbl94} and Colpi, Campana, \\& Treves (1993) \\markcite{cct93} for ONSs embedded in Giant Molecular Clouds (GMC) and by Zane \\etal (1996) \\markcite{zztt96} for ONSs in the solar proximity. However, despite the intense observational efforts, the search for ONSs produced, up to now, just a handful of candidates, out of which only two, RX J18653.5-3754 (Walter, Wolk, \\& Neuh\\\"auser 1996\\markcite{wwn96}) \\markcite{wwn96} and RX J0720.4-3125 (Haberl \\etal 1996\\markcite{hab96}, 1997\\markcite{hab97}), seem indeed promising. Moreover, recent analyses of ROSAT fields in the direction of GMCs (Belloni, Zampieri, \\& Campana 1996\\markcite{bzc96}; Motch \\etal 1997\\markcite{mo97}) placed rather stringent upper limits on the number of ONSs in ROSAT images. In all these cases, the predicted number of ONSs turns out to be too close to the observed number of non optically identified sources (NOIDs) which should comprise many other types of objects, like white dwarfs, active coronae and AGNs. In the field investigated by Motch {\\it et al.\\/}\\markcite{mo97}, for example, $\\sim 5$ sources should be detectable above 0.02 count$\\rm s^{-1}$ (Zane \\etal 1995\\markcite{zan95}) while the number of NOIDs, at the same limiting flux, is 8. Since these are sites where accretion should be the most effective, observations seem to suggest that current theoretical predictions are in excess of at least a factor $\\sim$ 5--10. Moreover, it was pointed out by Zane \\etal (1995)\\markcite{zan95} that most optimistic models produce a source density of $\\sim 10$ deg$^{-2}$ in ROSAT PSPC pointings at the limiting flux of $10^{-3}$ count $\\rm s^{-1}$, embarrassingly close to the average density of NOIDs, $\\sim 30$ deg$^{-2}$. A even more dramatic reduction of the number of ONSs could follow from the recent investigation of ROSAT sources in dark clouds and high latitude molecular clouds by Danner (1996)\\markcite{dan96}. For these fields predictions seem to be in excess by a factor $\\sim 10-100.$ Although the paucity of ONSs may be partly related to the heating of the ISM by the source itself (Blaes, Warren, \\& Madau 1995)\\markcite{bwm95}, it could be also closely connected to the properties of NSs at birth and/or to the long term evolution of their physical parameters. Present estimates on the observability of ONSs rely on a number of physical hypotheses concerning the velocity distribution of pulsars at birth and the evolution of the magnetic field (e.g. Treves, \\& Colpi 1991\\markcite{tc91}; Blaes, \\& Madau 1993\\markcite{bm93}; Colpi, Campana, \\& Treves 1993\\markcite{cct93}; Zane \\etal 1995\\markcite{zan95}). The velocity distribution of ONSs was derived from the evolution, over the lifetime of the Galaxy, of the pulsar velocity distribution of Narayan, \\& Ostriker (1990\\markcite{no90}, see also Paczy\\'nski 1990\\markcite{pa90}) while a relic magnetic field $\\lesssim 10^{10}$ G was assumed a priori. Disappointingly, both the velocity distribution of NSs and the long--term evolution of the magnetic field are still affected by large uncertainties which make them highly controversial issues. The actual number of detectable ONSs depends crucially on the number of low velocity neutron stars, which are those accreting at the highest rate. This, in turn, is directly related to the NS velocity distribution at birth. Lyne, \\& Lorimer (1994)\\markcite{ll94} and Hansen, \\& Phinney (1998)\\markcite{hp98} derived a pulsar velocity distribution containing fewer slow objects than in Narayan, \\& Ostriker (1990)\\markcite{no90}, but this result is in contrast with other recent investigations by Hartman (1997)\\markcite{ha197} and Hartman \\etal (1997)\\markcite{ha297}, which indicate that the number of pulsars with $v\\lesssim 40$ kms$^{-1}$ is the same as in Narayan \\& Ostriker or even higher. Observations show that neutron stars with magnetic fields $\\ll 10^{12}$ G, possibly $10^9$ G, are definitely present in LMXBs and in millisecond pulsars, suggesting that the magnetic field have decayed in these systems. Instead, no firm conclusions have been reached yet for isolated objects (see e.g. Srinivasan 1997\\markcite{sri97}; Wang 1997\\markcite{wa97}). Theoretical results are far from being univocal and predict either exponential/power--law field decay (Ostriker \\& Gunn 1969\\markcite{og69}; Sang \\& Chanmugam 1987\\markcite{sc87}; Goldreich \\& Reisenegger 1992\\markcite{gr92}; Urpin, Chanmugam \\& Sang 1994\\markcite{urp}; Miri 1996; Urpin, \\& Muslimov 1992) or little or no decay at all within the age of the Galaxy (Romani 1990\\markcite{ro90}; Srinivasan \\etal 1990\\markcite{sri90}; Goldreich \\& Reisenegger 1992\\markcite{gr92}; Pethick, \\& Sahrling 1995\\markcite{ps95}; see also Lamb 1991\\markcite{la91} for a revue). Statistical analyses based on observations of isolated radio pulsars give equally ambiguous results (Narayan \\& Ostriker 1990\\markcite{no90}; Sang \\& Chanmugam 1990\\markcite{sc90}; Bhattacharya \\etal 1992\\markcite{bha92}), owing in part to the difficulty of treating selection effects (Lamb 1992\\markcite{la92}). At the present, there is no clear evidence of field decay during the pulsar phase, but this does not preclude the possibility of field decay over longer timescales. Different approaches to pulsar statistics led, independently, to the conclusion that, if the magnetic field decays, then it probably does so over a timescale $\\sim 100$ Myr, well above the characteristic pulsar lifetime (Srinivasan 1997\\markcite{sri97}, Hartman \\etal 1997\\markcite{ha297}; Lorimer, Bailes, \\& Harrison 1997). Very recently, however, Wang (1997)\\markcite{wa97} and Konenkov, \\& Popov (1997)\\markcite{kp97} noticed that a substantial decay of the magnetic field should have occured precisely in the ONS candidate RX J0720.4-3125. This object may have nevertheless a different origin resulting from common envelope evolution in a binary system. In the last hypothesis the decay of the $B$--field may have been induced by accretion (see e.g. Shibazaki \\etal 1989\\markcite{shi89}; Romani 1990\\markcite{ro90}). Keeping in mind these uncertainties, the case of RX J0720.4-3125 may provide, for the first time, evidence for a decay of the magnetic field in aging, isolated NSs. If evolution leaves a relic field $\\sim 10^8-10^9$ G a large fraction of the entire NS population is expected to accrete the ISM, since major spin--down by propeller occured over the lifetime of the stars, as pointed out by Blaes, \\& Madau (1993)\\markcite{bm93}. Actually, previous investigations assumed that {\\it nearly all\\/} ONSs accrete from the ISM (see again Treves, \\& Colpi 1991; Blaes, \\& Madau 1993\\markcite{bm93}). At the low rates typical of ONSs embedded in the ISM, however, the conditions under which accretion is possible are rather stringent, since the inflowing gas may or may not penetrate inside the accretion and the Alfven radii according to the star velocity, spin period and magnetic field. As a consequence the number of ONSs accreting today may depend sensibly on the distribution of the stellar parameters at birth but, even more, on the long--term variation of the spin and magnetic field. The aim of this investigation is to explore the role of the magnetic field evolution in affecting theoretical predictions for ONSs observability. In order to do this, we relax the assumption of a constant, relic magnetic field and consider a model in which spin evolution causes the core field to migrate to the crust where dissipation processes (such as electron--phonon scatterings and scatterings on impurities) drive the field decay (see e.g. Srinivasan \\etal 1990\\markcite{sri90}; Urpin, \\& Muslimov 1992\\markcite{urp92}). Several evolutionary tracks have been computed and the properties of the ONS candidate RX J0720.4-3125 are easily recovered. We show that the paucity of ONSs detections can be accounted for if the magnetic field decays on a characteristic timescale $\\sim 10^8-10^9$ yr. The approach presented here suggests that, in close analogy with what is done for pulsars, the study of ONSs statistics may reveal a promising probe for constraining the spin and magnetic field evolution of neutron stars. ", + "conclusions": "In this paper we have shown that the expected number of ONSs observable with ROSAT can be sensibly reduced if the star magnetic field decays on timescales $\\sim 10^8-10^9$ yr. \\footnote{Soon after this paper was submitted, we became aware that a similar conclusion was reached by Livio, Xu, \\& Frank (1998)\\markcite{lxf98}.} This is primarily due to the lower fraction of NSs which are in the accretion phase now, but also to a softening of the spectra which may bring the peak of the emission outside ROSAT T--band if $\\tau_c=10^9$ yr. It follows that, if the paucity of detected ONSs is to be explained in terms of a global reduction in the number of accretion--powered sources, decay times $\\tau_c\\lesssim 10^8$ yr, should be invoked. For intermediate values of $\\tau_c$, $\\sim 10^9$ yr, the fraction of accreting ONSs is $\\sim 50-60\\%$, too large to match observational limits, but only $\\sim 15\\%$ of the sources falls in the ROSAT T--band, yielding again an acceptable number of {\\it detectable\\/} objects. This last figure should be regarded as indicative, being based on the assumption that all sources emit a blackbody spectrum from the polar caps. If the decay time is $\\sim 10^{10}$ yr or larger, spin--down by propeller is so effective that nearly all ONSs are accreting today. In this case the surface magnetic field is close to the initial value, $B\\sim 10^{12}$ G, and deviations from a pure blackbody spectrum are to be expected. Unfortunately, no detailed radiative transfer calculations are available for accretion atmospheres onto magnetized NSs, at least to our present knowledge. Some hints may be, nevertheless, derived from the analysis of radiative transfer in cooling neutron stars atmospheres (Shibanov \\etal 1992\\markcite{shi92}), since the input physics of the two model is quite similar. The emerging picture indicates that magnetic effects tend, somehow, to balance the hardening found by Zampieri \\etal (1995)\\markcite{ztzt95} in the case of unmagnetized accretion atmospheres, so that, for $B\\sim 10^{12}$ G, spectra are, globally, not too different from a blackbody emitted by the polar caps. With this proviso, figures in table 1 tell us that all accreting sources should be indeed within ROSAT T--bandpass and show that decay times $\\tau_c\\gtrsim 10^{10}$ yr provide a fraction of potentially detectable ONSs equal to that derived in previous investigations. Our main conclusion is that models of spin--induced magnetic field decay with a crustal dissipation time shorter than the age of the Galaxy are able to bring theoretical predictions of ONSs detectability closer to present observational limits. It is interesting to note that ohmic diffusion times between $10^{8.5}$ and $10^9$ yr can also explain the low values of $B$ observed in millisecond pulsars, which experienced spin induced field decay during binary evolution in wide low--mass systems (Miri, \\& Bhattacharya 1994\\markcite{mb94}). Although our present results seem indeed to point to a decay of magnetic field in ONSs, a firmer conclusion may be reached once the number of sources actually present in ROSAT fields is computed evolving the ONSs distribution function together with the spin and the magnetic field. Two other issues deserve further discussion. In computing the spin--down rate in the propeller phase the original expression of Illarionov, \\& Sunyaev (1975)\\markcite{is75} was used. It should be noticed that the physics governing the propeller is rather complex and still far from a complete understanding. It has been shown (see e.g. Ghosh 1995\\markcite{go95}) that the characteristic time for angular momentum losses strongly depends on the details of the accretion process and can be longer than what is assumed here. If the propeller is less efficient in braking the star, the number of accreting ONSs reduces further. On the other hand numerical simulations (Wang, \\& Robertson 1985\\markcite{wr85}) seem to indicate that propeller spin--down can go even faster that in equation (\\ref{propeller}). The second point concerns the field decay model. If the field is confined only in the stellar crust (see e.g. Urpin, \\& Muslimov 1992\\markcite{urp92}) the emerging picture may change since the magnetic field is not coupled to the period. In the model of Urpin, \\& Muslimov $\\tau_c$ is computed from the microphysics and varies in time, depending on the thermal evolution of the stellar crust. We have repeated the calculation using equation (\\ref{bsur}) with $B_c=0$ and $\\tau_c$ estimated from figure 5 of Urpin, \\& Muslimov, and found that the number of accreting ONSs again reduces. Even if the initial surface field is as large as $10^{13}$ G only $\\lesssim 1\\%$ of the NS population is the accretion phase now. On the opposite, an exponential decay of the surface field with $\\tau_c\\gtrsim 10^9$ yr, as inferred by Pethick, \\& Sahrling (1995) in the case of a non--vanishing core field, do not seem to solve the problem of the excess of potentially visible ONSs. The conclusion that many ONSs should be in the propeller stage seem rather robust if the magnetic field decays over a characteristic $\\tau_c\\sim 10^{9}$ yr, irrespective of the details of the model. Is there a way to observe directly these objects ? The power associated to the propeller derives ultimately from the rotational energy, and therefore one can deduce that it should be smaller than that released by accretion of matter onto the star surface typically by a factor $\\sim 100$ (see e.g. Treves, Colpi, \\& Lipunov 1993\\markcite{tcl93}). The accreted mass, however, could accumulate at the Alfven radius until the pressure required to overcome the centrifugal barrier is built up, giving raise to recurrent episodes of accretion. If this is the case, ONSs in the propeller phase could be observed as transient sources (Treves, Colpi, \\& Lipunov 1993\\markcite{tcl93}). More definite predictions about the observational appearance of these sources, including spectral properties, luminosity, recurrence times, demand, however, for a more realistic modelling of the physics governing the propeller." + }, + "9802/astro-ph9802204_arXiv.txt": { + "abstract": "Using high precision parallaxes from the Hipparcos catalog, we construct H-R diagrams for two samples of bright stars. The first is a magnitude-limited sample that is over $90\\%$ complete and uses uniform photometry from the Catalog of WBVR Magnitudes of Northern Sky Bright Stars ($\\delta > -14^\\circ$). This sample shows a smooth distribution of stars along the main sequence, with no detectable gaps. The second contains all of the stars closer than 100 parsecs in the Hipparcos catalog with $\\delta < -12^\\circ$. Uniform spectroscopy from the Michigan Spectral Survey shows that some stars which appear on the main sequence in the H-R diagram, particularly those in the $0.2 < B-V < 0.3$ region that has been labeled the B\\\"ohm-Vitense gap, are classified as giants by the MK system of spectral classification. Other gaps that have been identified in the main sequence are also affected by such classification criteria. This analysis casts doubt on the existence of the B\\\"ohm-Vitense gap, which is thought to result from the sudden onset of convection in stars. The standard identification of main sequence stars with luminosity class V, and giants with luminosity class III, must be reconsidered for some spectral types. The true nature of the stars that lie on the main sequence in the H-R diagram, but which do not have luminosity class V designations, remains to be investigated. ", + "introduction": " ", + "conclusions": "" + }, + "9802/astro-ph9802174_arXiv.txt": { + "abstract": "{The terrestrial fossil record shows that the exponential rise in biodiversity since the Precambrian period has been punctuated by large extinctions, at intervals of $40$ to $140$ Myr. These mass extinctions represent extremes over a background of smaller events and the natural process of species extinction. We point out that the non-terrestrial phenomena proposed to explain these events, such as boloidal impacts (a candidate for the end-Cretaceous extinction), and nearby supernovae, are collectively far more effective during the solar system's traversal of spiral arms. Using the best available data on the location and kinematics of the Galactic spiral structure (including distance scale and kinematic uncertainties), we present evidence that arm crossings provide a viable explanation for the timing of the large extinctions.} ", + "introduction": " ", + "conclusions": "" + }, + "9802/astro-ph9802342_arXiv.txt": { + "abstract": "We present the results of X-ray spectroscopy of a flux-limited sample of seven middle-aged supernova remnants (SNRs) in the Large Magellanic Cloud (LMC): N23, N49, N63A, DEM71, N132D, 0453$-$68.5, and N49B. We constructed self-consistent nonequilibrium ionization SNR models assuming a Sedov solution for the dynamical evolution, and then applied the resulting spectral models to the data obtained by the Solid-state Imaging Spectrometer onboard the {\\it Advanced Satellite for Cosmology and Astrophysics}. All the remnants were reasonably well described by the model, allowing us to derive accurate values for their physical parameters, i.e., ages, densities, initial explosion energies, and metal abundances. The derived explosion energies vary from $5\\times 10^{50}$ erg to $6\\times 10^{51}$ erg. A restricted subset of the sample exists for which the ionization and Sedov dynamical ages agree quite well under the assumption that the electron and ion temperatures are not fully equilibrated at the shock front; for these four SNRs the mean value of the initial explosion energy is $(1.1\\pm 0.5)\\times 10^{51}$ erg. We show that it is likely that the other three remnants exploded within preexisting cavities in the interstellar medium. The limits on high energy X-ray emission ($\\gsim$3 keV) that we present indicate that these SNRs do not contain very luminous pulsar-powered synchrotron nebulae, in general agreement with our picture of them as evolved, middle-aged remnants. We find statistical evidence for enrichment by supernova ejecta in the sense that smaller remnants show a somewhat higher mean metallicity than the larger ones. In the case of DEM71, the putative remnant of a Type Ia supernova, the derived abundance of iron is about a factor of two larger than the other remnants in the sample. These things being said, however, the derived abundances are in general dominated by swept-up interstellar material and so we use the SNR sample to estimate the mean LMC gas-phase abundances. We find that the astrophysically common elements from oxygen to iron are less abundant than the solar values by factors of 2--4. Overall these results are consistent with previous ones based on optical and UV data, but do not show the anomalous overabundance of magnesium and silicon seen by others. ", + "introduction": "Studies of the supernova remnants (SNRs) in the Large Magellanic Cloud (LMC) are essential for elucidating the details of SNR evolution, nucleosynthesis, the nature and environments of supernova progenitors, and so on. As probes of the interstellar medium (ISM), SNRs also yield information on the energy balance and chemical composition of their environments, which serves as input to larger questions of galactic chemical evolution and the star formation history of the Cloud. The relative closeness of the LMC plus its well-determined distance (we use 50 kpc throughout) means that accurate physical quantities can be derived for individual remnants in the galaxy. Also, due to the generally low interstellar absorption toward the LMC, it is possible to detect X-rays in the important 0.5--2.5 keV energy band, which includes emission lines from some of the most abundant metals in the Universe, such as K-shell transitions of highly-ionized atoms of oxygen, neon, magnesium, silicon and sulfur and L-shell lines of ionized iron. The Solid-state Imaging Spectrometer (SIS) onboard the {\\it Advanced Satellite for Cosmology and Astrophysics} (\\asca), with its superior energy resolution compared to previous broadband X-ray spectrometers, gives us the first real opportunity for studying these issues quantitatively in the X-ray regime. Mathewson et al.\\ (1983, 1984, 1985) identified 32 SNRs in the LMC on the basis of optical, radio, and X-ray observations. The remnants range in size from 2 pc to about 100 pc, and therefore span the range of evolutionary phases, from young ejecta-dominated remnants, through middle-aged adiabatic remnants, to old remnants in the radiative phase. At least 12 of the LMC remnants are bright enough X-ray emitters ($F_{X} > 10^{-11}$ erg s$^{-1}$ cm$^{-2}$ in the 0.15--4.5 keV band), that we can perform detailed spectroscopic study with \\asca. First results have already been reported from \\asca\\ on three small LMC remnants, N103B, 0509$-$67.5, and 0519$-$69.0 (Hughes et al.\\ 1995), that show strong emission lines originating from ejecta characteristic of thermonuclear supernovae, i.e., Type Ia SNe. For the research described herein, we observed seven middle-aged remnants, N23, N49, N63A, DEM71, N132D, 0453$-$68.5, and N49B, with \\asca\\ and performed detailed spectral analysis of their SIS data. This is all the LMC SNRs larger than 5 pc (radius) and brighter than $10^{11}$ ergs s$^{-1}$ cm$^{-2}$ over the soft keV X-ray band, which includes the three brightest thermal X-ray--emitting remnants in the LMC (N132D, N63A, and N49), plus four less luminous ones that cover a range of evolutionary states. Previous X-ray spectroscopic studies of N49, N63A, N132D, and N49B were reported by Clark et al.\\ (1982) based on Solid State Spectrometer (SSS) data from the \\EO. The SSS data clearly showed weak emission lines of Mg, Si, and S indicative of thermal emission, although the analysis, which used somewhat unphysical two-temperature collisional ionization equilibrium plasma models, left some doubt about the reliability of the derived elemental abundance values. More recently Hwang et al.\\ (1993) analyzed all the available \\einstein\\ data on N132D using a phenomenological nonequilibrium ionization (NEI) model and derived abundances of the heavy elements that were considerably lower than the solar values. Both analyses suggest that the bulk of the X-ray emission from these remnants comes from swept-up ISM rather than supernova ejecta and, as we show below, our \\asca\\ data confirm this finding. This observational fact gives us confidence that the SNR abundances we measure indeed correspond to those of the ISM in the LMC. The elemental abundances of the Magellanic Clouds are of great importance to extragalactic astronomy. As the nearest major galaxies to our own, they are well-studied and serve as the stepping stones for studies of more distant systems. In the current chemical composition of the Clouds we see at least some of the material returned to the ISM by stars of all types through winds and supernovae and, using models of stellar evolution and nucleosynthesis, we can constrain the history of star formation in these galaxies. It is well known that the Clouds have undergone quite different star formation histories from the Galaxy, based on their differing chemical properties (see, for example, Tsujimoto et al.~1995). In the past 20 years with the advent of high quality optical and UV spectra, the study of the chemical abundances of the LMC has made great progress. The earlier results of this era (summarized in Dufour 1984) showed that the metallicity of LMC \\hii\\ regions was significantly less than solar; the observed oxygen abundance, for example, was only about \\onethird\\ the solar value. More recent studies (Luck \\& Lambert 1992, Russell \\& Dopita 1992, hereafter RD92) have, in general, tended to confirm these findings using optical spectra and detailed modeling of \\hii\\ regions, SNRs, Cepheids, and supergiant stars in the LMC. However, even though considerable effort has been expended in this area of research over the past decade, there remains considerable room for improvement. In particular, certain critical elements, such as Mg and Si, are not easily studied in the optical and UV bands due to a paucity of appropriate atomic transitions. The work we present here represents the first systematic attempt to measure the gas-phase LMC abundances in a new, independent manner using X-ray spectroscopy of SNRs. Our spectral analysis also allows us, for the first time, to derive accurate values for the temperature, ionization timescale, and intrinsic X-ray emissivity of the hot gas in LMC SNRs. When combined with their physical radii, which are accurately known because of the well-established distance to the LMC, we have four independent observational constraints on the dynamical evolution. Only three of these quantities are necessary to fully specify the evolutionary state according to the Sedov (1959) similarity solution, so in fact the observational data result in an overconstrained system. This allows us to go beyond merely assuming the Sedov solution and instead allows to actually establish whether or not the Sedov model is a correct representation for the evolution of the SNRs in our sample. In the following we present our \\asca\\ analysis of the selected LMC SNR sample. In \\S 2 we briefly present the observational details. In contrast to the phenomenological models used in previous X-ray studies of LMC SNRs, we have developed realistic self-consistent models that incorporate the dynamical evolution of remnants. Section 3 provides a description of the spectral model we developed and compares our calculations to similar published ones. Results of the Sedov model fits, as well as considerations of the applicability of the model, are given in \\S 4. Discussions of the derived abundances are in \\S 5. In \\S 6 we examine the high energy X-ray emission ($\\gsim$3 keV) of the remnant sample. We comment on individual remnants in \\S 7 and summarize in the final section. ", + "conclusions": "We have carried out a systematic analysis of the \\asca\\ SIS X-ray spectra of seven luminous SNRs in the LMC. Our work has resulted in the following conclusions. \\begin{description} \\item[(1)] The spectral data are described well by a NEI model that includes the time evolution of density and temperature given by the Sedov similarity solution for the dynamical evolution of SNRs. This model was a better description of the data than any single or multiple component phenomenological NEI model that we tried. \\item[(2)] The X-ray data provide 4 independent constraints on the 3 independent parameters of the Sedov model. The observational constraints are: radius, shock temperature, emission measure, and ionization timescale. The parameters of the Sedov model are age, initial explosion energy, and ambient ISM density. There is an additional ambiguity in the models that arises from the unknown relationship between the energy imparted to electrons and ions at the shock front. We model two extreme situations: full equilibration (in which the electron and ion temperatures are assumed to be equal) and Coulomb equilibration (in which the electrons gain energy from the ions though Coulomb collisions only). \\item[(3)] Four of the remnants in the sample, N23, N49, DEM71, and 0453$-$68.5, are fully consistent with Sedov evolution under the assumption of Coulomb equilibration. For this group of remnants the mean explosion energy is $(1.1\\pm 0.5)\\times 10^{51}$ ergs in agreement with the canonical value. The key to discriminating between the full and Coulomb equilibration models comes from comparing the Sedov dynamical age with the age determined from the ionization timescale. \\item[(4)] When the derived parameters of the other three SNRs, N63A, N132D, and N49B, are examined in detail, they appear to be internally inconsistent. In particular the age determined by the ionization timescale is considerably less than the Sedov dynamical age, regardless of whether one assumes the full or Coulomb equilibration models. Furthermore the initial SN explosion energies are large $\\gsim$$3\\times 10^{51}$ ergs. Both of these discrepancies can be explained by invoking a scenario in which the remnants exploded within preexisting cavities or bubbles in the ISM and that the X-ray emission we see now comes from the blast wave interacting with the dense material at the cavity wall. This was previously suggested for N132D to explain the discrepancy between this remnant's Sedov age and the kinematic age determined from the expansion of high-velocity oxygen-rich filaments. Our independent estimate of N132D's age from its ionization timescale confirms this earlier result and it gives confidence that N63A and N49B are actually two more examples of this phenomenon. This result is not particularly surprising since it is precisely the manner in which massive stars are predicted to modify their environment. \\item[(5)] We find statistical evidence for enrichment by supernova ejecta in the sense that smaller remnants show a somewhat higher mean metallicity than the larger ones. On the other hand, the mean metallicity does not correlate with swept-up mass. In the case of the Balmer-dominated SNR DEM71, which is likely to be the remnant of a Type Ia supernova, the derived abundance of iron is about a factor of two larger than the other remnants in the sample. This corresponds to an excess of iron over the amount in the ISM of $\\sim$0.06 $M_\\odot$, which is about 10\\% of the total amount of iron ejected by a Type Ia SN. \\item[(6)] All things being considered, however, the middle-aged, evolved SNRs in our sample are in general dominated by swept-up ISM and so can be used to estimate the mean LMC gas-phase abundances. We find that the common elements from oxygen to iron have abundances 0.2--0.4 times solar, consistent with previous results based on optical and UV data, but without the anomalous overabundance of Mg and Si seen by others. This work demonstrates the validity of using X-ray spectra of SNRs to measure the current chemical composition of interstellar gas and it has the potential to provide a significant new constraint on the chemical evolution and star formation history of the Cloud. \\item[(7)] We set limits on the hard ($\\gsim$3 keV) X-ray flux of the remnants that correspond to luminosities of from $1.9\\times 10^{35}\\rm\\, ergs\\, s^{-1}$ to $2.2\\times 10^{36}\\rm\\, ergs\\, s^{-1}$ in the 0.2--4.0 keV band. These limits indicate that the remnants in the sample do not contain very luminous pulsar-powered synchrotron nebulae, as is consistent with our picture of them as evolved middle-aged SNRs. However, the quoted luminosity range is consistent with the possiblity that they contain weaker synchrotron nebulae with pulsars losing energy at a slower rate than is the case in the Crab Nebula, for example. The discovery of such pulsar-powered synchrotron nebulae would definitively identify the progenitor as a massive star that underwent a core collapse SN. \\end{description}" + }, + "9802/astro-ph9802048_arXiv.txt": { + "abstract": "s {We present preliminary results on galactic Dark Matter (DM), halo structure, and galaxy evolution. We show how during the first Gyr of the evolution of a $10^{10} M_{\\odot}$ dwarf elliptica feed-back from stars (SN\\ae~ and stellar winds) leads to an extended constant density isothermal core with radius of 0.15 the virial radius $R_{200}$. We also present first results on galaxy merging as a possibile scenario to form ellipticals, studying in particular how the details of the merging evolution vary as a function of the mass ratio of the interacting galaxies.} \\normalsize\\baselineskip=15pt ", + "introduction": "Numerical N-body and gasdynamical simulations have become a major tool to investigate how galaxies formed and evolved. Massively parallel computers, like Cray T3D or T3E supercomputers, have allowed to run high resolution simulations of large scale structure formation (VIRGO or GIF projects), leading to impressive results on the distribution of DM in different cosmological scenarios (Jenkins et al. 1997). On galaxy scales, Navarro et al. (1996a) showed that galactic DM halos should follow a universal density profile. In other words, the violent, collisionless dynamical relaxation processes during the formation of DM halos lead to equilibrium profiles with similar shapes, independent of the halo mass, the initial density fluctuation spectrum, and the adopted cosmological model. However, recent studies (Persic et al. 1996; Burkert \\& Silk 1997) have found evidence, in real galaxies, of significant departures from this universal profile both in the inner and in the outer galactic regions. \\begin{figure*}[top] \\vbox{\\vsize=10cm\\vskip-2.1cm% \\centerline{\\psfig{file=fig1.ps,height=12cm,width=12cm}} \\caption{Time evolution of the DM density profile.}} \\end{figure*} In particular, the presence of extended DM cores seems to be a firm result. The occurence and the development of cores with costant density is not yet completely understood. It has been suggested that the discrepancy between observations and Cold Dark Matter (CDM) predictions could be solved by assuming secular processes in the baryonic component which may also affect the innermost halo regions. An analytic approach used by Navarro et al. (1996b) has suggested that such processes might produce core only in low mass galaxies. In this contribution we present a fully Nbody simulation of the formation of a dwarf elliptical galaxy to investigate whether baryonic feedback can actually be the cause of the cores developments. Tothis aim we use our own TreeSPH code (for details see Carraro et al 1997 and Lia et al 1998). ", + "conclusions": "This paper is a report on two ongoing projects in Galaxy Formation. Within the context of the monolithic collapse scenario we have proposed a possible explanation for the formation of a DM core in a dwarf elliptical, emphasising the role of secular processes related to stellar feedback. In the context of the merging scenario for the formation of ellipticals, we have presented preliminary results on the merging of two spiral-like galaxies for two cases with different mass ratios." + }, + "9802/astro-ph9802081_arXiv.txt": { + "abstract": "Using numerical techniques we studied the global stability of cooling flows in giant elliptical galaxies. As an initial equilibrium state we choose the hydrostatic gas recycling model \\cite{kritsuk96}. Non-equilibrium radiative cooling, stellar mass loss, heating by type Ia supernovae, distributed mass deposition, and thermal conductivity are included. Although the recycling model reproduces the basic X-ray observables, it appears to be unstable with respect to the development of inflow or outflow. In spherically symmetry the inflows are subject to a central cooling catastrophe, while the outflows saturate in a form of a subsonic galactic wind. Two-dimensional axisymmetric random velocity perturbations of the equilibrium model trigger the onset of a cooling catastrophe, which develops in an essentially non-spherical way. The simulations show a patchy pattern of mass deposition and the formation of hollow gas jets, which penetrate through the outflow down to the galaxy core. The X-ray observables of such a hybrid gas flow mimic those of the equilibrium recycling model, but the gas temperature exhibits a central depression. The mass deposition rate $\\dot M$ consists of two contributions of similar size: (i) a hydrostatic one resembling that of the equilibrium model, and (ii) a dynamical one which is related to the jets and is more concentrated to the centre. For a model galaxy, like NGC~4472, our 2D simulations predict $\\dot M\\approx2$~M$_{\\odot}$~yr$^{-1}$ within the cooling radius for the advanced non-linear stage of the instability. We discuss the implications of these results to H$\\alpha$ nebulae and star formation in cooling flow galaxies and emphasize the need for high-resolution 3D simulations. ", + "introduction": "X-ray observations of normal early-type galaxies with the {\\em Einstein} observatory have revealed an extended diffuse thermal emission from the hot (0.5-2~keV) component of their interstellar medium [ISM] \\cite{forman....79}. With luminosities $L_{\\rmn{X}}\\sim10^{39}$-$10^{42}$~erg~s$^{-1}$ this hot gas contributes $10^9-10^{10}$~M$_{\\odot}$ to the total galaxy mass in its optical confines (Forman, Jones \\& Tucker 1985\\nocite{forman..85}). The X-ray luminosities correlate with the optical (blue) luminosities of early-type galaxies, although with a large spread of $L_{\\rmn{X}}$ at a fixed $L_{\\rmn{B}}$ [Canizares, Fabbiano \\& Trinchieri \\shortcite{canizares..87}; Donnelly, Faber \\& O'Connell \\shortcite{donnelly..90}; Eskridge, Fabbiano \\& Kim \\shortcite{eskridge..95b}]. For nearby bright ellipticals the surface brightness distributions in X-rays and in the optical are similar as far as the interaction of the galaxy X-ray emitting gas with the surrounding medium is unimportant (Trinchieri, Fabbiano \\& Canizares 1986\\nocite{trinchieri..86}). {\\em ROSAT} observations revealed correlations of the temperature of the hot gas with the stellar velocity dispersion and with the abundances for a complete sample of 43 elliptical galaxies \\cite{davis.96}. Galaxies with higher velocity dispersions tend to have higher, approximately solar, abundances and higher diffuse gas temperatures. In all cases the gas is substantially (a factor of $\\sim2$) hotter than the kinetic temperature of the luminous stars indicating the presence of dark haloes. The analysis of spectral properties for a sample of 12 early-type galaxies observed with {\\em ASCA} has revealed systematically lower abundances with a mean value of about 0.3 solar and has confirmed the relationship between stellar velocity dispersion and gas temperature \\cite{matsumoto.....97}. Gas temperature profiles for elliptical galaxies determined with {\\em ROSAT} and {\\em ASCA} have been found to be surprisingly uniform for projected radii $r\\le10\\:r_{\\rmn{e}}$ ($r_{\\rmn{e}}$ is the so-called effective or half-light radius). They exhibit a positive temperature gradient out to $\\sim3\\:r_{\\rmn{e}}$ followed by a leveling off or gradual decrease toward larger radii \\cite{brighenti.97}. In addition to hot X-ray emitting gas early-type galaxies have been shown to have a cooler component ($10^7-10^8$~M$_{\\odot}$) detectable at 100~$\\mu$m \\cite{jura...87} as well as warm ionized gas ($10^3-10^5$~M$_{\\odot}$ in the central 1-10~kpc) producing optical emission lines \\cite{caldwell84,demoulin-ulrich..84,phillips...86,buson........93,singh...95,macchetto......96}. Initially, steady-state cooling flow models were invoked to interpret X-ray observations \\cite{thomas....86,sarazin.87,sarazin.88,vedder..88,sarazin.89}. This, in essence, empirical approach was based on the simple idea that if the cooling time-scale is shorter than the Hubble time near the centre, but still longer than the gravitational time-scale, then a very subsonic quasi-hydrostatic inflow {\\em must} take place. However, the global stability of these early models has not been proved yet. Later time-dependent hydrodynamic modeling has demonstrated that, in general, such accretion flows are unstable and suffer cooling catastrophes at their centres, which were difficult to reconcile with observations \\cite{meiksin88,murray.92}. Secular variations of the stellar mass-loss rate and supernova heating rate due to stellar evolution in the galaxy allowed for a description of the evolution of the hot ISM in ellipticals on the Hubble time-scale [D'Ercole et al. \\shortcite{d'ercole...89}; Ciotti et al. \\shortcite{ciotti...91}; David, Forman \\& Jones (1990, 1991) \\nocite{david..90,david..91a}]. The resulting models evolved through up to three consecutive evolutionary stages: the wind, outflow and inflow phases. Accordingly, most of present-day ellipticals are in the outflow phase, while a few of the brightest galaxies may have already experienced their transition to the inflow regime. Although the problem of the cooling catastrophe was not resolved, an important outcome of spherically symmetric evolutionary modeling was the understanding that the use of steady-state cooling flow models to diagnose mass accretion rates in dynamical situations can lead to erroneous results \\cite{ciotti...91,murray.92}. In the absence of direct measurements of the flow velocity a systematic study of global hydrodynamic and thermal stability properties of the hot ISM in ellipticals can provide a selection criterion for realistic flow regimes. Starting from the essential physics of cooling flows, we have constructed a hydrostatic equilibrium model, which describes a stable {\\em in situ} recycling of gas shed by stars through local thermal instabilities into a condensate, which can become the material for further star formation \\cite{kritsuk96}. In this paper we probe the global stability of the recycling model in the inflow-outflow context. Section 2 describes the basic assumptions of the hot ISM model and of the underlying galaxy model. Section 3 gives the details of the numerical method used in the simulations. The evolution of spherically symmetric perturbations and their stabilization by a conductive heat flux are considered in Section 4. In Section 5 we abandon the restrictive assumption of spherical symmetry and study the development of instabilities by means of two-dimensional simulations. This allows for a better insight into the complex physics of cooling flows. The discussion and conclusions can be found in Sections 6 and 7, respectively. ", + "conclusions": "Using numerical technique we probed the stability of the hydrostatic equilibrium recycling model for hot gaseous coronae of giant elliptical galaxies with respect to a variety of perturbations. In the absence of heat conduction the equilibrium appears to be unstable on a time scale of the order of the thermal time for the gas at the galactic centre. The physical reason for this instability is rooted in the properties of the source terms. The energy and mass sinks due to local thermal instabilities essentially modify the elastic properties of the gas. A slow contraction of a gas element amplifies the condensation and cooling in it and, thereby, the mass inflow into the compressed region. As a result, the density rises, providing further growth of losses and accelerating the contraction in a runaway regime. A slow expansion of the gas element instead decreases the losses and increases the specific heating per unit mass, which further decreases the losses and accelerates the element expansion. The expansion and heating saturate as the gas temperature approaces the characteristic temperature of the heat source. Both types of unstable behaviour can be identified in our simulations. Spherically symmetric perturbations trigger either gas inflow, which ends up with a cooling catastrophe in the core of a galaxy, or a subsonic outflow regime, which efficiently removes the gas from the central region. In the first case the X-ray surface brightness at the centre of the galaxy considerably exceeds observed values, while in the second case the hot gas is practically invisible in X-rays. Thermal conduction is able to maintain a stable equilibrium only against low-amplitude perturbations. Spherically symmetric, global density perturbations with amplitudes $\\ge10$~\\% remain unstable. In case of inflows the conductive heat flux cannot prevent the cooling catastrophe at the centre, but can delay it in time considerably. On the other hand, when the thermal conductivity is not suppressed, the gas velocities of outflows saturate at much lower values. This gives rise to quasi-steady-state solutions with an X-ray luminosity only slightly lower than that of the equilibrium model. Using axisymmetric perturbations, we were able to study the 2D hydrodynamics of the cooling catastrophe for the first time. In axial symmetry the system exhibits an qualitatively different behaviour. A set of narrow cooling gas streams, flowing towards the centre through a global subsonic galactic wind, develops in response to random velocity perturbations of the equilibrium recycling model. In the strongly non-linear regime the characteristic averaged X-ray observables of such hybrid gas flows (e.g., the surface brightness and the temperature profiles) mimic the characteristics of the initial equilibrium state, while the equilibrium itself appears to be considerably violated. These results indicate the need for a three-dimensional time-dependent treatment of the problem, which would be able to reveal non-linear instability saturation mechanisms. At the same time, they cast doubts on using steady-state `cooling flow' type models, based on the assumption of quasi-hydrostatic equilibrium, for recovering mass deposition rates in the central galactic regions." + } +} \ No newline at end of file